{
    "9804/astro-ph9804073_arXiv.txt": {
        "abstract": "We introduce a statistical quantity, known as the $K$ function, related to the integral of the two--point correlation function. It gives us straightforward information about the scale where clustering dominates and the scale at which homogeneity is reached. We evaluate the correlation dimension, $D_2$, as the local slope of the log--log plot of the $K$ function. We apply this statistic to several stochastic point fields, to three numerical simulations describing the distribution of clusters and finally to real galaxy redshift surveys. Four different galaxy catalogues have been analysed using this technique: the Center for Astrophysics I, the Perseus--Pisces redshift surveys (these two lying in our local neighbourhood), the Stromlo--APM and the 1.2 Jy {\\it IRAS} redshift surveys (these two encompassing a larger volume). In all cases, this cumulant quantity shows the fingerprint of the transition to homogeneity. The reliability of the estimates is clearly demonstrated by the results from controllable point sets, such as the segment Cox processes. In the cluster distribution models, as well as in the real galaxy catalogues, we never see long plateaus when plotting $D_2$ as a function of the scale, leaving no hope for unbounded fractal distributions. ",
        "introduction": "The standard cosmology is based on the assumption that the Universe must be homogeneous on very large scales. Several pieces of evidence support this assumption: the homogeneity and isotropy of the microwave background radiation \\cite{cobe} and some aspects of the large scale distribution of matter \\cite{peb89} seem to strongly advocate uniformity on scales bigger than about $200 \\, h^{-1}$ Mpc (where $H_0= 100 h$ km s$^{-1}$ Mpc$^{-1}$).  However the presence of very large features in the galaxy distribution like the Bootes void \\cite{kir81} or the Great Wall \\cite{gel89} which span a scale length of the order of $100 \\, h^{-1}$ Mpc calls the actual scale of homogeneity into question. Moreover other authors consider the assumption of homogeneity just a theoretical prejudice not necessarily supported by the observational evidence quoted above. They defend the alternative idea of an unbounded fractal cosmology \\cite{col92}. Guzzo (1997) argues against this interpretation on the basis of a careful handling of the data.  The spatial two--point correlation function is the statistical tool mainly used to describe the clustering in the Universe (Peebles 1980, 1993). However, because of the integral constraint \\cite{p80}, one cannot estimate it at very large distances from the currently available redshift surveys. In order to study clustering in the regime where it is not very strong, we have only two possibilities: either we extend the size of the redshift catalogues or we use  alternative statistical descriptors. The approach described in this paper points in the latter direction. In the same line, other authors \\cite{fis93,par94,tad96} have tried to measure the power--spectrum on large scales directly from galaxy catalogues. Einasto \\& Gramman (1993) studied the transition to homogeneity by means of the power--spectrum and found a relation between the correlation transition scale and the spectral transition scale (turnover in $P(k)$).  We introduce the quantity called $K(r)$, which is related to the correlation function $\\xi(r)$. The novelty of our approach lies essentially in the fact that we shall use a cumulant quantity instead of a  differential quantity such as $\\xi(r)$. Although for a point process the functions $\\xi(r)$ and $K(r)$ are well defined, what we measure from the galaxy catalogues are just estimators of those functions. One of our main claims is that the estimators for $K(r)$ are more reliable than the most currently used estimators for $\\xi(r)$ and that makes its use recommendable (especially in three-dimensional processes and at large scales) despite its somewhat less informative character. ",
        "conclusions": "We should like also to comment briefly on the relation of $K$ with the correlation function $\\xi(r)$. Both play their role in the analysis of the point pattern and, as Stoyan \\& Stoyan (1996) say, their relation is similar to that between the distribution function and the probability density function in classical statistics. The use of a cumulative quantity such as $K$ avoids binning in distance, which is often a source of arbitrariness for $\\xi$ \\cite{rip92}. Let us explain why $\\xi$ does suffer from the hindrance of splitting the information into disjoint bins. When one estimates $\\xi(r)$ in $[r,r+dr]$, it is assumed that within that bin the correlation function is constant, and since this is obviously not true, the larger the bin the larger the error, but we cannot make arbitrarily small the size $dr$ of the bin, because in that case we would not find any pairs. In other words, $\\xi(r)$ has an additional source of bias, not present in $K$, due to the smoothing caused by averaging over pairs of points close to but not exactly $r$ units apart of each other (Stein 1996).  The correlation length ($r_0| \\xi(r_0)=1$) is just the scale at which the density of galaxies is, on average, twice the mean number density. At smaller scales the pair correlations are due to non--linear perturbations, but homogeneity is not reached till $\\xi(r_{\\rm \\tiny hom}) \\sim 0$. The main interest of $K$ is that it permits us to study clustering precisely in that \\lq\\lq difficult'' range where $r_0 < r < r_{\\rm \\tiny hom}$, which cannot be reached by $\\xi$ because in this range the errors on the estimates of $\\xi$ are comparable with their values, while the difference $K-K_{\\rm \\tiny Pois}$ is still meaningful.  As a concluding remark, we want to stress that an unbiased estimator of a quantity related with the correlation integral, known as the $K$ function, has been applied to cosmological simulations and galaxy samples. This function, extensively used in the field of spatial statistics, provides a nice measure of clustering. The border correction used here does not waste any data points and does not introduce spurious homogeneization, giving reliability to the evaluation of this function at large scales. Through the slope of $K$ we are able to calculate $D_2$, which is an indicator of a possible fractal behaviour of the point process at a given scale range. The clear physical meaning of $K$ and $D_2$ helps us easily interpret the clustering properties of different models of structure formation at different scales.  Regarding the analysis of the galaxy redshift surveys, we have seen that the estimator of the $K$ function is robust in the sense that it does not depend on the shape of the study region and provides us with reliable information about the point patterns over a wide range of scales. The behaviour of the local dimension $D_2$ for the real galaxy samples is particularly interesting to proponents of various fractal models of large--scale structure. If a constancy of $D_2$ with the scale is a necessary condition for having a fractal point pattern (although it should not be sufficient as we have seen with the Cox process [see also Stoyan (1994) for more examples]), it is a neat conclusion of our analysis that the galaxy distribution does not even hold the necessary condition. The analysis presented here will provide a conclusive test to discover the scale at which the distribution of the matter in the Universe is really homogeneous when applied, in the near future, to the bigger and deeper galaxy catalogues which will be soon ready for common use.     \\subsection*{ACKNOWLEDGEMENTS}   This work has been partially supported by an EC Human Capital and Mobility network (contract ERB CHRX-CT93-0129) and by the Spanish DGES project n. PB96-0707. We thank prof. Stoyan for bringing the Cox model to our attention and for useful conversations and comments. We thank R. Croft, S. Paredes, R. Trasarti--Battistoni  and R. van de Weygaert  for kindly allowing us to use their samples and programs, as well as T. Buchert, J. Schmalzing, M. Stein and specially M. Kerscher for very interesting discussions and comments. The authors want to thank the anonymous referee for his/her valuable comments and suggestions."
    },
    "9804/gr-qc9804034_arXiv.txt": {
        "abstract": "Primordial black holes may form in the early Universe, for example from the collapse of large amplitude density perturbations predicted in some inflationary models. Light black holes undergo Hawking evaporation, the energy injection from which is constrained both at the epoch of nucleosynthesis and at the present. The failure as yet to unambiguously detect primordial black holes places important constraints. In this article, we are particularly concerned with the dependence of these constraints on the model for the complete cosmological history, from the time of formation to the present. Black holes presently give the strongest constraint on the spectral index $n$ of density perturbations, though this constraint does require $n$ to be constant over a very wide range of scales. ",
        "introduction": "Black holes are tenacious objects, and any which form in the very early Universe are able to survive until the present, unless their Hawking evaporation is important. The lifetime of an evaporating black hole is given by \\begin{equation} \\frac{\\tau}{10^{17} \\, {\\rm sec}} \\simeq \\left( \\frac{M}{10^{15} \\, {\\rm grams}} \\right)^3 \\,. \\end{equation} From this we learn that a black hole of initial mass $M \\sim 10^{15}$g will evaporate at the present epoch, while for significantly heavier black holes Hawking evaporation is negligible. Another mass worthy of consideration is $M \\sim 10^{9}$g, which leads to evaporation around the time of nucleosynthesis, which is well enough understood to tolerate only modest interference from black hole evaporation by-products.  Several mechanisms have been proposed which might lead to black hole formation; the simplest is collapse from large-amplitude, short-wavelength density perturbations. They will form with approximately the horizon mass, which in a radiation-dominated era is given by \\begin{equation} \\label{hormass} M_{{\\rm HOR}} \\simeq 10^{18} \\, {\\rm g} \\, \\left( \\frac{10^7 \\, {\\rm GeV}}{T} \\right)^2 \\,, \\end{equation} where $T$ is the ambient temperature. This tells us that any black holes for which evaporation is important must have formed during very early stages of the Universe's evolution. In particular, formation corresponds to times much earlier than nucleosynthesis (energy scale of abut $1\\,$MeV), which is the earliest time that we have any secure knowledge concerning the evolution of the Universe. Any modelling of the evolution of the Universe before one second is speculative, and especially above the electro-weak symmetry breaking scale (about $100 \\,$GeV) many possibilities exist. Note also that although we believe we understand the relevant physics up to the electro-weak scale, the cosmology between that scale and nucleosynthesis could still be modified, say by some massive but long-lived particle. In this article we will consider the standard cosmology and two alternatives \\cite{gl,glr}.  We define the mass fraction of black holes as \\begin{equation} \\beta \\equiv \\frac{\\rho_{{\\rm pbh}}}{\\rho_{{\\rm tot}}} \\,, \\end{equation} and will use subscript `i' to denote the initial values. In fact, we will normally prefer to use \\begin{equation} \\alpha \\equiv \\frac{\\rho_{{\\rm pbh}}}{\\rho_{{\\rm tot}}-\\rho_{{\\rm pbh}}} = \\frac{\\beta}{1-\\beta} \\,, \\end{equation} which is the ratio of the black hole energy density to the energy density of everything else. Black holes typically offer very strong constraints because after formation the black hole energy density redshifts away as non-relativistic matter (apart from the extra losses through evaporation). In the standard cosmology the Universe is radiation dominated at these times, and so the energy density in black holes grows, relative to the total, proportionally to the scale factor $a$. As interesting black holes form so early, this factor can be extremely large, and so typically the initial black hole mass fraction is constrained to be very small.  The constraints on evaporating black holes are well known, and we summarize them in Table~\\ref{massfrac}. This table shows the allowed mass fractions at the time of evaporation. An additional, optional, constraint can be imposed if one imagines that black hole evaporation leaves a relic particle, as these relics must then not over-dominate the mass density of the present Universe \\cite{BCL:PBH}. For black holes massive enough to have negligible evaporation, the mass density constraint is the only important one (though in certain mass ranges there are also microlensing limits which are somewhat stronger).  \\begin{table}[t] \\caption[massfrac]{\\label{massfrac} Limits on the mass fraction of black holes at evaporation.} \\begin{tabular}{|c|c|c|} \\hline \\hline Constraint & Range & Reason \\\\ \\hline $\\alpha_{\\rm{evap}} < 0.04$ & $10^{9}$ g $< M < 10^{13}$ g & Entropy per baryon\\\\ & & at nucleosynthesis \\cite{var:ent} \\\\ \\hline $\\alpha_{\\rm{evap}} < 10^{-26} \\frac{M}{m_{{\\rm Pl}}}$ & $M \\simeq 5\\times10^{14}$~g & $\\gamma$ rays from current\\\\ & & explosions \\cite{var:gam} \\\\ \\hline $\\alpha_{\\rm{evap}} < 6\\times10^{-10} \\left( \\frac{M}{m_{{\\rm Pl}}}\\right)^{1/2}$ & $10^{9}$~g $ < M <10^{11}$~g & n$\\bar{\\rm{n}}$ production \\\\ & & at nucleosynthesis \\cite{var:neu} \\\\ \\hline $\\alpha_{\\rm{evap}} < 5\\times10^{-29} \\left( \\frac{M}{m_{{\\rm Pl}}}\\right)^{3/2}$ & $10^{10}$~g $< M < 10^{11}$~g & Deuterium destruction \\cite{lin:deu} \\\\ \\hline $\\alpha_{\\rm{evap}} < 1\\times10^{-59}\\left( \\frac{M}{m_{{\\rm Pl}}}\\right)^{7/2}$ & $10^{11}$~g $< M < 10^{13}$~g & Helium-4 spallation \\cite{var:he4}\\\\ \\hline \\end{tabular} \\vspace*{2pc} \\end{table}  We will study three different cosmological histories in this paper, all of which are currently observationally viable. The first, which we call the standard cosmology, is the minimal scenario. It begins at some early time with cosmological inflation, which is necessary in order to produce the density perturbations which will later collapse to form black holes. Inflation ends, through the preheating/reheating transition (which we will take to be short), giving way to a period of radiation domination. Radiation domination is essential when the Universe is one second old, in order for successful nucleosynthesis to proceed. Finally, radiation domination gives way to matter domination, at a redshift $z_{{\\rm eq}} = 24\\,000\\,\\Omega_0 h^2$ where $\\Omega_0$ and $h$ have their usual meanings, to give our present Universe.  The two modified scenarios replace part of the long radiation-dominated era between the end of inflation and nucleosynthesis. The first possibility is that there is a brief second period of inflation, known as thermal inflation \\cite{ls}. Such a period is unable to generate significant new density perturbations, but may be desirable in helping to alleviate some relic abundance problems not solved by the usual period of high-energy inflation. The second possibility is a period of matter-domination brought on by a long-lived massive particle, whose eventual decays restore radiation domination before nucleosynthesis. For definiteness, we shall take the long-lived particles to be the moduli fields of superstring theory, though the results apply for any non-relativistic decaying particle. ",
        "conclusions": "Although black hole constraints are an established part of modern cosmology, they are sensitive to the entire cosmological evolution. In the standard cosmology, a power-law spectrum is constrained to $n < 1.25$, presently the strongest observational constraint on $n$ from any source. Alternative cosmological histories can weaken this to $n < 1.30$, and worst-case non-gaussianity \\cite{BP} can weaken this by another 0.05 or so, though hybrid models giving constant $n$ give gaussian perturbations. Finally, we note that while the impact of the cosmological history on the density perturbation constraint is quite modest due to the exponential dependence of the formation rate, the change can be much more significant for other formation mechanisms, such as cosmic strings where the black hole formation rate is a power-law of the mass per unit length $G\\mu$. After all, the permitted initial mass density of black holes does increase by many orders of magnitude in these alternative cosmological models."
    },
    "9804/astro-ph9804245_arXiv.txt": {
        "abstract": "We describe a method for the extraction of spectra from high dispersion objective prism plates. Our method is a catalogue driven plate solution approach, making use of the Right Ascension and Declination coordinates for the target objects.  In contrast to existing methods of photographic plate reduction, we digitize the entire plate and extract spectra off-line. This approach has the advantages that it can be applied to CCD objective prism images, and spectra can be re-extracted (or additional spectra extracted) without having to re-scan the plate. After a brief initial interactive period, the subsequent reduction procedure is completely automatic, resulting in fully-reduced, wavelength justified spectra.  We also discuss a method of removing stellar continua using a combination of non-linear filtering algorithms.  The method described is used to extract over 12,000 spectra from a set of 92 objective prism plates. These spectra are used in an associated project to develop automated spectral classifiers based on neural networks. ",
        "introduction": "The MK classification of stellar spectra (Morgan, Keenan \\& Kellman 1943\\nocite{morgan_43a}) has been an important tool in the workshop of stellar and galactic astronomers for more than a century. While improvements in astrophysical hardware have enabled the rapid observation of digital spectra, our ability to efficiently analyze and classify spectra has not kept pace. Traditional visual classification methods are clearly not feasible for large spectral surveys. In response to this, we have been working on a project to develop automated spectral classifiers (von Hippel et~al.\\ 1994; Bailer-Jones 1996; Bailer-Jones et~al.\\ 1997, 1998).  These classifiers, which are based on supervised artificial neural networks, can rapidly classify large numbers of digital spectra.  The development of these classification techniques has required a large, representative set of previously classified spectra. The most suitable data has been the spectra from the Michigan Spectral Survey (Houk 1994)\\nocite{houk_94a} and the accompanying MK spectral type and luminosity class classifications listed in the {\\it Michigan Henry Draper} (MHD) catalogue (Houk \\& Cowley 1975; Houk 1978, 1982; Houk \\& Smith-Moore 1988). \\nocite{houk_75a}\\nocite{houk_78a}\\nocite{houk_82a}\\nocite{houk_88a} This paper describes the data reduction techniques we developed to extract and process these spectra. ",
        "conclusions": "\\begin{figure} \\centerline{ \\psfig{figure=fig9.eps,width=0.5\\textwidth,angle=0} } \\caption{Distribution of spectral types for each luminosity class. The dotted line represent giants (III), the dashed line subgiants (IV) and the solid line dwarfs (V).} \\label{dist_B} \\end{figure}  This paper has described a method for extracting spectra from objective prism images.  The method has been developed for the reduction of a set of photographic objective prism plates, but because the spectral extraction and processing takes place entirely in software using the complete digitized plate, it can equally well be applied to CCD objective prism images.  The extraction process is driven by a set of catalogue Right Ascension and Declination positions, so a direct image of each field is not required. After an initial interactive period taking one or two minutes, the subsequent reduction is automatic, taking approximately one hour on a modest-sized SUN Sparc IPX to process a single plate (i.e.\\ extract about 150 spectra).  The reduction method described in this paper has been used to extract a set of over 12,000 high-quality spectra.  From this, a subset of over 5,000 normal spectra was selected which had reliable two-dimensional (spectral type and luminosity class) classifications listed in the MHD catalogue.  The frequency distribution of the various stellar classes in this set is shown in Figure~\\ref{dist_B}. This data set is used in accompanying papers to produce automated systems for classifying and physically parametrizing stellar spectra (Bailer-Jones et~al.\\ 1997, 1998).  In the interests of extending spectral classification to more distant stellar populations, spectra of stars fainter than B $\\sim 12$ are required. This could be achieved with a CCD objective prism survey. Although the technique described can only extract objects with known Right Ascension and Declination coordinates, the HST Guide Star Catalogue (e.g.\\ Lasker et al.\\ 1990)\\nocite{lasker_90a}, which lists 19 million objects brighter than 16$^{th}$ magnitude, could be used as a driver for extraction. However, Bailer-Jones (unpublished, 1996) has also modified the method to extract unwidened spectra from CCD objective prism images in the absence of any coordinates, using an algorithm to locate local flux peaks.  The method can be applied to spectra at different spectral resolutions and wavelength coverages, provided a suitable line exists for the second plate solution."
    },
    "9804/astro-ph9804135_arXiv.txt": {
        "abstract": "We have undertaken near-continuous monitoring of the Seyfert 1 galaxy NGC 7469 in the X-ray with \\xte\\ over a $\\sim 30$~d baseline. The source shows strong variability with a root-mean-square (rms) amplitude of $\\sim 16$~per cent, and peak--to--peak variations of a factor of order 2.  Simultaneous data over this period were obtained in the ultraviolet (UV) using \\iue, making this the most intensive X-ray UV/X-ray variability campaign performed for any active galaxy.  Comparison of the continuum light curves reveals very similar amplitudes of variability, but different variability characteristics, with the X--rays showing much more rapid variations.  The data are not strongly correlated at zero lag. The largest absolute value of the correlation coefficient occurs for an anticorrelation between the two bands, with the X-ray variations leading the UV by $\\sim 4$~d.  The largest positive correlation is for the ultraviolet to lead the X-rays by $\\sim 4$~d. Neither option appears to be compatible with any simple interband transfer function.  The peak positive correlation at $\\sim 4$d occurs because the more prominent peaks in the UV light curve appear to lead those in the X-rays by this amount. However, the minima of the light curves are near-simultaneous.  These observations provide new constraints on theoretical models of the central regions of active galactic nuclei.  Models in which the observed UV emission is produced solely by re-radiation of absorber X-rays are ruled out by our data, as are those in which the X-rays are produced solely by Compton upscattering of the observed UV component by a constant distribution of particles.  New or more complex models must be sought to explain the data.  We require at least two variability mechanisms, which have no simple relationship.  We briefly explore means by which these observations could be reconciled with theoretical models. ",
        "introduction": "\\label{Sec:Introduction}  The origin of the continuum emission of Active Galactic Nuclei (AGN) -- which covers an extremely broad band -- is not well understood.  In a number of high luminosity sources, it appears that this emission peaks in the ultraviolet (UV), the so-called ``Big Blue Bump'' (Shields 1978; Malkan \\& Sargent 1982).  Strong, and apparently non-thermal X-ray flux is also a persistent property of AGN (e.g., Marshall \\etal\\ 1981). The X-ray emission covers a wide band from at least 0.1-100~keV, and can be described by a power-law form (Mushotzky, Done \\& Pounds 1993).  The Big Blue Bump is often identified as the thermal output of an accretion disk (henceforth the accretion disk model e.g., Shakura \\& Sunyaev 1973). Heat is generated by viscous dissipation in the disk, which then radiates in the optical/UV regime for black hole masses typical of AGN (e.g., Sun \\& Malkan 1989).  An alternative origin for the UV continuum emissions has been suggested by both observational and theoretical considerations. Guilbert \\& Rees (1988) postulated that the UV need not be internally generated in the accretion disk, but could arise via absorption and thermal re-emission (hereafter referred to as ``thermal reprocessing'') of X-rays in optically thick gas close to the central engine.  The material - which could be but does not necessarily have to be the disk -- would imprint features on the X-ray spectra (e.g., Lightman \\& White 1988; George \\& Fabian 1991; Matt, Perola \\& Piro 1991). Such features have been found (e.g., Nandra \\& Pounds 1994) and suggest that approximately half of the incident X-rays are absorbed in the optically thick material. Spectroscopic observations of strong gravitational and Doppler effects in the iron K$\\alpha$ line profiles of Seyfert galaxies (e.g., Tanaka \\etal\\ 1995; Nandra \\etal\\ 1997) suggest that this material lies extremely close to the central black hole and is probably in the form of a disk (Fabian \\etal\\ 1995).  The bulk of the continuum photons absorbed in the gas should then be re-emitted at the characteristic thermal temperature of the material. For dense gas close to the central engine, and particularly for  ``standard'' accretion disks, this should be in the optical/UV.  Most models for the X--ray continuum of AGN are based on the idea that lower-energy photons are Compton scattered by a population of hot electrons and/or pairs (which we refer to as ``upscattering'' models e.g., Sunyaev \\& Titarchuk 1980; Svensson 1983; Guilbert, Fabian \\& Rees 1983). The seed photons are often assumed to be those in the blue bump.  Specific models differ primarily in their assumptions about, e.g., the geometry of the system (e.g. Haardt \\& Maraschi 1991, 1993; Haardt, Maraschi \\& Ghisellini 1994; Stern \\etal\\ 1995), the question of whether the electron population has a thermal or non-thermal distribution, the importance of pairs (e.g., Zdziarski \\etal\\ 1990, 1994).  These models have been successful in explaining various observations. The goal underlying the exploration of the models is the discovery of the process responsible for the generation of the copious energy output of AGN.  While the case for accreting supermassive black holes is becoming compelling, the method by which the rest-mass energy of the material is converted into radiative energy is still highly uncertain. Some specific questions which remain about the emission mechanisms include:  \\begin{enumerate} \\item{How important is viscous dissipation in the generation of the UV?} \\item{What proportion of the UV arises via thermal reprocessing of X--rays?} \\item{What is the seed population for upscattering into the X--rays?} \\item{What mechanism accelerates the particles which up-scatter these seed photons?} \\end{enumerate}  A powerful way of investigating these questions is by variability campaigns. These have already reaped rich rewards in the study of AGN emission lines via ``reverberation mapping'' (e.g. Peterson 1993, Netzer \\& Peterson 1997, and references therein). These emission line campaigns, however, also had strong implications for the generation of the continua, which we shall discuss below. The models discussed above all imply strong connections between the continuum emission in different bands.  For example, the accretion disk emission could cover an extremely broad band, depending on the temperature profile of the disk. The thermal reprocessing model predicts that the X-rays should be generating UV emission. The upscattering model suggests the converse. By observing the variability in these bands, therefore, we can make inferences as to which of the various processes is in operation and to what degree. In particular, simultaneous X-ray/UV data should be the most revealing.  A number of AGN have been monitored simultaneously at optical/UV and X-ray energies.  Leaving aside blazars, the best-studied sources are NGC~4151, NGC~5548 and NGC 4051. In the first two objects, there is evidence for a correlation between the two bands. The best-sampled (and therefore most reliable) case is NGC~4151, in which the 1455~\\AA\\ and 2--10~keV flux appears to correlate well on all time scales from hours to a year (Perola \\etal\\ 1986; Edelson \\etal\\ 1996).  In NGC~5548, the flux in the two bands is also well correlated on time scales from days to 1 year (Clavel \\etal\\ 1992). In both sources however, the correlation appears to break down during one very large UV outburst.  NGC~4051 shows different behavior, in that the X-ray emission showed large-amplitude (factor $\\sim 2$) variability, while the optical emission remained steady to within a few per cent when observed over a $\\sim 2$~d baseline (Done \\etal\\ 1990).  For completeness, we also mention the results obtained for other non-blazar AGN, though their significance is marginal due to the small number of simultaneous observations and/or the short duration of the campaigns.  In Fairall~9, the slow decline of the 2--10~keV flux mimics the secular fading of the UV and optical continuum from 1978 to 1985 (Morini \\etal\\ 1986).  The UV-optical versus X-ray flux correlation seems to hold in NGC~4593 (Santos-Ll\\'{e}o \\etal\\ 1995), whereas in MCG-8-11-11 (Treves \\etal\\ 1990), 3C120 (Maraschi \\etal\\ 1991) and 3C~273 (Courvoisier \\etal\\ 1990) the two wavebands appear to be independent of each other.  These previous attempts at determining the relationship between the components have obviously left some ambiguity. This is perhaps not surprising as generally the sampling of the light curves has been rather poor.  In order to provide an improved dataset, a campaign of near-continuous \\xte\\ and \\iue\\ monitoring of NGC 7469 was undertaken over a $\\sim 1$~month baseline. The results of the campaign in terms of the relationship of the X-ray and UV variability are the subject of this paper. We have effectively divided the paper into two halves. Sections 2-4 discuss the observational results exclusively, which are then summarized in Section 5. Section 6 then investigates the implications of the observational results within the framework of the models discussed above and suggest possible ways of reconciling the data with models. ",
        "conclusions": "We have investigated the relationship between the X-ray and UV emission in NGC 7469 on time scales of hours-weeks.  The poor correlation between the X-ray and UV light curves at zero lag may be considered a surprising result because, as mentioned in \\S1, some previous experiments suggested a good correlation, and little if any time lag between the variations in the two bands (e.g. NGC 5548, NGC 4151).  No other AGN, however, has been monitored as intensively and for such a long duration.  Variability information has important implications for the physical mechanisms responsible for the production of the X-ray and UV emission in AGN. If the flux changes in two bands are correlated, this suggests some causal link between them. A time lag between the bands then shows which component drives the other. If X-ray variations lead those in the UV, this is strongly suggestive of thermal reprocessing. If the opposite is observed, this strongly favors upscattering. With a simple transfer function this delay should be similar for all ``events'' in the light curve. No such simple behavior was observed during our campaign and the interpretation is less straightforward. Our data require modifications to the simplest ideas about the emission mechanisms in the UV and X-ray.  \\subsection{Model ingredients}  It is widely accepted that the radiative energy from AGN originates as the rest-mass energy of accreting material. The conversion process must include a mechanism to accelerate the particles which produce the X-rays as they are the highest-energy photons to carry an appreciable fraction of the luminosity.  In principle, the remainder of the AGN spectrum could be produced by thermal reprocessing following absorption of some fraction of this X-ray continuum.  Viscous dissipation in the accretion flow, however, could dominate the observed UV emission.  In many scenarios there is also a radiative connection between the UV and X-ray emission regions: X-rays can be produced by upscattering of UV seed photons and UV emission can be produced by re-radiation of absorbed X-rays. In these circumstances, a substantial number of factors can affect the observed variability:  \\begin{enumerate} \\item Changes in the physical properties (e.g. optical depth, temperature, geometry) of the particle distribution which produces the X-rays \\item Instabilities in the accretion flow \\item The geometry and size of the X-ray and UV emission regions \\item Anisotropy of the radiation fields, which might include relativistic effects close to the black hole \\item The temperature distribution of the absorbing medium \\item The importance of feedback, in which variations in each band affect the other \\item Changes in occulting/absorbing media \\end{enumerate}  \\subsection{Implications of the NGC 7469 data}  Our data have a number of implications for the processes which produce the UV--to--X-ray emissions of NGC 7469 and the observed variability, which we now discuss. First we consider our observations in the context of historical data and the spectral energy distribution of NGC 7469.  \\subsubsection{The broad-band perspective}  The mean X-ray and UV fluxes observed for NGC 7469 during our campaign are very much typical of the respective historical means for this source. The mean flux in the 2--10 keV band, based on historical observations over the period 1979-1993, is $\\sim 3 \\times 10^{-11}$\\,erg cm$^{-2}$ s$^{-1}$. Furthermore, the range of historical variability is very similar to that observed during our campaign. This suggests that we sampled a large fraction of source variability in NGC 7469 during our one-month campaign, although we observe no obvious flattening of the PDS (Fig.~\\ref{fig:pds}) at the lowest frequencies.  Chapman, Geller and Huchra (1985) derived a mean value of $4.6 \\times 10^{-14}$\\,erg~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$ for the 1430--1460~\\AA\\ continuum flux of NGC~7469 from 10 \\iue\\ observations in 1979--1982. Similarly, Edelson, Pike and Krolik (1990), reported a mean value of $4.8 \\times 10^{-14}$\\,erg~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$ for the continuum flux at 1450~\\AA\\ (rest wavelength) from 16 \\iue\\ observations in 1979--1985. With a mean UV flux at 1485~\\AA\\ (observed wavelength) of $4.0 \\times 10^{-14}$\\,erg~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$ (W97), NGC~7469 was thus neither particularly faint nor exceptionally bright during our campaign.  The optical--to--X-ray spectral index of NGC~7469, $\\alpha_{ox}$~=~1.22, is not significantly different from that of, e.g., NGC~5548 ($\\alpha_{ox}$~=~1.25), or the mean index for Seyfert galaxies (Kriss, Canizares and Ricker 1980).  The spectral energy distribution of NGC 7469 is shown in Fig.~\\ref{fig:sed}. The only real peculiarity of NGC~7469 is the presence of a circumnuclear starburst ring within $1.\\!''5$ of its nucleus. Genzel \\etal\\ (1995) estimate that the starburst accounts for two-thirds of the source bolometric luminosity and it may dominate the IR emission.  However, it only contributes 4 percent to the observed X-ray flux (Perez-Olea and Colina 1996). The starburst should be invariant on the time scales sampled here and have no effect on the X--ray/UV variability.  \\subsubsection{The X--ray continuum}  The presence of rapid variations in the X-rays which are not seen in the UV implies that the particle distribution responsible for the X-rays is variable.  The X-ray emission mechanism is highly uncertain, but as mentioned in the introduction, many models have concentrated on Compton upscattering of seed UV photons by a population of hot electrons and/or pairs (e.g. Haardt \\& Maraschi 1991, 1993). Our observations show that if the observed UV photons are the seed population, then the rapid variations of the X-rays do not arise from variations in the seed.  Either the optical depth, temperature or geometry of the upscattering region must be changing.  In the latter case, changes in the distribution of active regions in the X-ray source, or kinematic effects can produce variability (e.g., Abramowicz \\etal\\ 1991; Haardt, Maraschi \\& Ghisellini 1997).  Longer-term variations are also observed in the X-ray flux and these could also arise from processes such as those just mentioned. The fact that the power spectrum shows no obvious features or break is supportive of this interpretation. The fluctuations have a similar amplitude to those in the UV and this suggests a connection between the bands.  This is intriguing. One possibility is that the long time-scale variability in the X--rays is due to changes in the UV seed population. In the simplest such interpretation - where the upscattering region is point-like and lies in the line of sight to a point-like seed source we expect a 1:1 correlation between the two bands with no lag. This is ruled out by our data, although we note that the minima appear to be very close in time. In more complex geometries, there may be time lags which will be in the sense that the X-ray variations follow those in the UV. The delay of $\\sim 4$~d between the peaks is superficial evidence in favor of upscattering. In that model, however, we expect the lag to be very short, being dependent primarily on the light travel time between the regions, modified by geometrical factors. A lag of $\\sim 4$~d, the only plausible conclusion here, seems rather long to be associated with these processes. In addition, we are unable to envisage a purely geometrical modification of a linear process which accounts for {\\it both} the relationship between the maxima and that of the minima.  \\subsubsection{The UV continuum}  We now consider the alternative that the X-ray emission drives the UV in the thermal reprocessing scenario.  The luminosities in the X-ray and UV bands are similar (Fig.~\\ref{fig:sed}) and thermal reprocessing can therefore be energetically important.  As stated above, the 2-10 keV observed flux of NGC 7469 at our epoch was $3.4\\times 10^{-11}$\\,erg cm$^{-2}$ s$^{-1}$. However, the X-ray emission of NGC 7469 covers a far wider band than this, with significant emission being observed down to $\\sim 0.1$~keV with ROSAT (Brandt \\etal\\ 1993) and most likely up to at least 100~keV as seen by OSSE (Zdziarski \\etal\\ 1995; Gondek \\etal\\ 1996). Estimates of the underlying photon index of the continuum are in the range $1.9-2.0$ after accounting for the effects of Compton reflection (Piro \\etal\\ 1990; Nandra \\& Pounds 1994).  We estimate the mean X--ray luminosity of NGC 7469 at our epoch to be $1.8\\times 10^{44}$\\,erg s$^{-1}$ in the 0.1-100~keV band. In the canonical thermal--reprocessing scenario about half of this luminosity should be absorbed in the accretion disk or other material. After estimating that fraction which is Compton scattered rather than absorbed (George \\& Fabian 1991) we conclude that the presence of the iron emission line and reflection hump in NGC 7469 indicate that a luminosity of $\\sim8\\times 10^{43}$\\,erg s$^{-1}$ of the X-ray emission of NGC 7469 is reprocessed and re-emerges as thermal emission.  Let us now suppose that all of this luminosity emerges in a single black body (which represents the narrowest physically-realistic spectrum) peaking close to 1315\\AA\\ ($kT \\sim 2$~eV).  Such a blackbody is almost sufficient to account for the continuum at 1315\\AA\\ (Figure~\\ref{fig:sed}).  Therefore, it is energetically possible that reprocessed X-rays produce some of the observed UV continuum and its variations. As in the case of upscattering, the simplest thermal reprocessing models predict a strong positive correlation between the bands, with any time lags being in the sense that the X-ray variations lead those in the UV. No such lag is observed. We do find that the strongest {\\it anti}--correlation of the datasets occurs for the X-rays leading the UV, but the interpretation of such a result is far from obvious and we do not comment on it further.  Even if aliasing has caused us to ``miss'' a positive correlation with a long ($\\sim 14$~d) lag between the UV and X--rays, any lag longer than a day or so is very difficult to explain. Any single transfer function which related the two bands would smooth the light curve of the responding band and reduce its amplitude. We observe, however, that the amplitudes of variability on long time scales are very similar. We therefore reject such a possibility. In this case it is even more difficult to envisage a complex geometry which can reproduce the light curves.  It therefore seems highly unlikely that any substantial proportion of the 1315\\AA\\ continuum of NGC 7469 arises from thermal reprocessing unless, for example, there is substantial anisotropy of the X-ray emission.  W97 found that the variations at longer UV wavelengths followed those at shorter wavelengths, but with a time lag of a fraction of a day.  Collier \\etal\\ (1998) have demonstrated that this trend continues into the optical, and Peterson \\etal\\ (1998) find this trend to be significant at no less than the 97~per cent confidence level. This, together with the rapidity of the variations in NGC 7469, is most easily explicable in terms of the thermal reprocessing hypothesis. However, our data essentially rule out models in which all the observed optical/UV flux is re-radiated X-ray emission. The optical/UV variability therefore requires either intra-band reprocessing, which is difficult from an energetics standpoint, or some other model.  Should we therefore conclude that the UV/optical continuum in NGC 7469 arises from direct emission by the accretion disk?  Our data offer no direct constraints on accretion disk models, as no explicit relationship between the X--ray and UV emissions is predicted by those models. Nonetheless, the fact that thermal reprocessing is strongly disfavored by our data has profound implications for the disk models. The rapid and wavelength-coherent variations in the optical and UV flux of AGN is difficult to reconcile with a standard $\\alpha$-disk (e.g. Krolik \\etal\\ 1991; Molendi, Maraschi \\& Stella 1992). Prior to our observations, it was conceivable that thermal reprocessing was responsible for these rapid variations.  There is now a clear need for a revision of accretion disk theory to account for these wavelength-independent variations without resorting to reprocessing.  \\subsubsection{The extreme ultraviolet (EUV) continuum and UV emission lines}  Given the presence of a typical iron K$\\alpha$ line and reflection hump in this source (Piro \\etal\\ 1990; Nandra \\& Pounds 1994) we are left with the question of where the putative reprocessed X-ray flux is emitted. One possibility is that the thermal reprocessing occurs in a molecular torus (Ghisellini, Haardt \\& Matt 1994; Krolik, Madau \\& Zycki 1994), in which case it might emerge in the infrared. Such an hypothesis would predict a narrow iron K$\\alpha$ line in the X-ray spectrum, whereas in many Seyfert galaxies these lines are extremely broad. The case of NGC 7469 is unclear, with Guainazzi \\etal\\ (1994) finding no evidence for a broad component and Nandra \\etal\\ (1997) finding marginal evidence. A conclusive determination requires a longer exposure with \\asca, but it seems highly likely that the iron K$\\alpha$ line and Compton hump in Seyfert 1 galaxies in general are produced extremely close to the central black hole (e.g., Nandra \\etal\\ 1997). We would therefore expect the reprocessed emission to emerge at a higher energy. As shown above, however, the thermally reprocessed X-rays only make a strong contribution to the observed optical and UV wavebands if the emission is strongly peaked at those wavelengths. It seems more likely that the emission covers a range of temperatures, in which case the reprocessed flux would be difficult to detect when spread over a wide band.  Alternatively, it could peak in the (unobserved) EUV band. Figure~\\ref{fig:sed} shows that this can indeed be the case. A blackbody of luminosity $8\\times 10^{43}$\\,erg s$^{-1}$ contributes less than 5~per cent of the flux at 1315\\,\\AA\\ as long as $kT>12$~eV. Intensity variations in such a component would be undetectable with \\iue. Similarly, the \\asca\\ spectrum constrains $kT<60$~eV. If the spectral form is broader than a single blackbody, the range of allowed temperatures is correspondingly wider. Interestingly, Brandt \\etal\\ (1993) reported evidence for a soft excess in the \\ros/PSPC data which can be modeled as a blackbody of $kT\\sim 110$~eV and a luminosity of $10^{43}$\\,erg s$^{-1}$.  This can be identified with the high energy tail of such a broad, reprocessed component.  The major UV emission lines are excited by unobservable EUV photons. An extrapolation of the X-ray spectrum observed by \\asca\\ (George \\etal\\ 1998) and of the UV spectrum into that band indicate roughly equal contributions at energies at which the lines are excited.  With the two components being poorly correlated at zero lag, it is therefore difficult to determine which will be the dominant EUV component at any given time.  We have suggested above that there may even be a third contributor to the EUV, the reprocessed X-rays.  In other words, the shape of the ionizing continuum changes with time. This effect could account for certain difficulties which have been encountered in explaining the emission line responses to the observed UV continuum in reverberation mapping experiments. Our observations suggest that the unseen EUV continuum is not directly related to the observed \\iue\\ flux, which therefore cannot be assumed to be a perfect representation of the continuum driving the line emission. It is also interesting to note that the emission-line light curves show long term trends which are not apparent in the continuum bands. This is most clearly demonstrated by Fig.~\\ref{fig:renorm_lc}, which shows the X-ray and UV continuum light curves, together with those of the Ly$\\alpha$ and C{\\sc iv} emission lines. These have all been renormalized to the $F_{\\rm var}$ value and thus the y-axes crudely represent the number of standard deviations from the mean. Both emission lines clearly show a long-term reduction in their flux which is not seen in either continuum.  \\subsection{Steps towards a new model}  In the light of the above, it is clear that new or more complex models must be sought to explain the data which have been obtained thus far, and particularly those described in this paper. Here we suggest some ways in which our new data might be reconciled with the existing paradigm by modification or extension. We emphasize that such a discussion is incomplete and {\\it ad-hoc}.  As we have stated above, it seems most likely that the X--ray flux which is absorbed when the iron K$\\alpha$ lines is being generated emerges in a relatively weak, broad component, that may peak the EUV/soft X-ray band. The emission in this band may well provide the crucial connection between the higher and lower-energy components.  A reasonable interpretation of the longer-timescale variability observed in our light curve is that the UV emission leads that in the X-rays, but with a variable lag. This suggests the dominant source of variations is in the seed population of an upscattering model. We do, however, bear in mind the caveat that the particle distribution of the upscattering medium must also be variable, to produce the most rapid variations. To explain the ``variable'' time lag, we suggest that there are multiple ``seed'' populations, which dominate at different times. In particular we suggest that the main source of 1315\\AA\\ photons is located at a distance of $\\sim 4$~lt d from the X--ray source and they are the dominant seed population when the source is in a high flux state, thus introducing a 4d ``lag'' between the X-rays and UV. When the 1315\\AA\\ flux is observed to decline, however, this allows other emitting regions to dominate the seed distribution. In particular we suggest at these times that EUV/soft X-ray photons are the dominant seed population. They arise from closer in and are therefore observed to have little or no lag with the X-rays.  As might be apparent from the above discussion, the primary X--rays, reprocessed X-rays and primary UV might well exist in a rather fine balance in the typical AGN. Future observational data on other objects of similar quality to that presented here, and preferably including the far-UV and soft X-ray bands, will be necessary for further progress and to establish the generality of the phenomena explored here."
    },
    "9804/astro-ph9804303_arXiv.txt": {
        "abstract": "We report a tentative detection with the IRAM 30m telescope of the LiH molecule in absorption in front of the lensed quasar B0218+357. We have searched for the $J = 0 \\rightarrow 1$ rotational line of lithium hydride at 444 GHz (redshifted to 263 GHz). The line, if detected, is optically thin, very narrow, and corresponds to a column density of N(LiH) = 1.6 10$^{12}$ cm$^{-2}$ for an assumed excitation temperature of 15 K, or a relative abundance LiH/H$_2 \\sim$ 3 10$^{-12}$. We discuss the implications of this result. ",
        "introduction": "Primordial molecules are thought to play a fundamental role in the early Universe, when stellar nucleosynthesis has not yet enriched the interstellar medium. After the decoupling of matter and radiation, the molecular radiative processes, and the formation of H$_2$, HD and LiH contribute significantly to the thermal evolution of the medium (e.g. Puy et al 1993, Haiman, Rees \\& Loeb 1996). Even at the present time, it would be essential to detect such primordial molecules, to trace H$_2$ in the low-metallicity regions (e.g. Pfenniger \\& Combes 1994, Combes \\& Pfenniger 1997). Unfortunately, the first transition of HD is at very high frequency (2.7 THz), and the first LiH line, although only at 444 GHz, is not accessible from the ground at $z=0$ due to H$_2$O atmospheric absorption. This has to wait the launching of a submillimeter satellite.  Although the Li abundance is low (10$^{-10}$-10$^{-9}$), the observation of the LiH molecule in the cold interstellar medium looks promising, because it has a large dipole moment, $\\mu = 5.9$ Debye (Lawrence et al.~1963), and the first rotational level is at $\\approx 21\\,\\rm K$ above the ground level, the corresponding wavelength is $0.67\\,\\rm mm$ (Pearson \\& Gordy 1969; Rothstein 1969). The line frequencies in the submillimeter and far-infrared domain have been recently determined with high precision in the laboratory (Plummer et al 1984, Bellini et al 1994). Because of the great astrophysical interest of this molecule (e.g. Puy et al 1993), an attempt has been made to detect LiH at very high redshifts ($z \\sim 200$) with the IRAM 30m telescope (de Bernardis et al 1993). It has been proposed that the LiH molecules could smooth the primary CBR (Cosmic Background Radiation) anisotropies, due to resonant scattering, or create secondary anisotropies, and they could be the best way to detect primordial clouds as they turn-around from expansion (Maoli et al 1996, but see also Stancil et al 1996, Bougleux \\& Galli 1997).  There has recently been some controversy about the abundance of LiH. The computations of Lepp \\& Shull (1984) estimated the LiH/H$_2$ abundance ratio in primordial diffuse clouds to be as high as 10$^{-6.5}$. With H$_2$/H $\\sim$ 10$^{-6}$, the primordial LiH/H ratio is $\\sim$10$^{-12.5}$. More recently, Stancil et al. (1996) computed an LiH/H abundance of $< 10^{-15}$ in the postrecombination epoch, since quantum mechanical computations now predict the rate coefficient for LiH formation through radiative association to be 3 orders of magnitude smaller than previously thought from semi-classical methods (Dalgarno et al 1996). In very dense clouds, however, three-body association reactions must be  taken into account, and a significant fraction of all lithium will turn into molecules. Complete conversion due to this process requires gas densities of the order $\\sim 10^9$\\,cm$^{-3}$, rarely found in the general ISM. However, taken other processes into account, such as dust grain formation, an upper limit to the LiH abundance is the complete conversion of all Li into molecular form, with LiH/H$_2$ $\\la 10^{-10}-10^{-9}$. With a LiH column density of 10$^{12}$\\,cm$^{-2}$, or N(H$_2$)$= 10^{22}$\\,cm$^{-2}$, the optical depth of the LiH line will reach $\\sim$1, in cold clouds of velocity dispersion of $2\\,\\rm km\\,s^{-1}$. The line should then be easily detectable in dense dark clouds in the present interstellar medium (like Orion where the column density reaches 10$^{23}$-10$^{24}$ cm$^{-2}$). This is a fundamental step to understand the LiH molecule formation, in order to interpret future results on primordial clouds, although the primordial abundance of Li could be increased by about a factor 10 in stellar nucleosynthesis (e.g. Reeves 1994). Once the Li abundance is known as a function of redshift, it could be possible to derive its true primordial abundance, a key factor to test Big Bang nucleosynthesis (either homogeneous or not).  Up to now, due to atmospheric opacity, no astrophysical LiH line has been detected, and the abundance of LiH in the ISM is unknown. The atmosphere would allow to detect the isotopic molecule LiD (its fundamental rotational line is at 251 GHz), but it has not been seen because of the low D/H ratio, and the expected insufficient optical depth of LiH \\footnote{the LiD line at 251 GHz is not covered in the 247-263 GHz survey of Orion by Blake et al 1986, but was observed at the McDonald 5m-telescope, Texas, see Lovas 1992; we have ourselves checked with the SEST telescope that no line is detected towards Sagittarius-B2 at this frequency. The 3$\\sigma$ upper limit to the LiD abundance towards SgrB2 is $1 \\times 10^{11}$\\,cm$^{-2}$.}.  Another method to avoid atmospheric absorption lines is to observe a remote object, for which the lines are redshifted into an atmospheric window. Here we report about the first absorption search for a LiH line at high redshift: the latter allows us to overcome the earth atmosphere opacity, and thanks to the absorption technique we benefit of an excellent spatial resolution, equal to the angular size of the B0218+357 quasar core, of the order of 1 milli-arcsec (Patnaik et al 1995). At the distance of the absorber (redshift $z=0.68466$, giving an angular size distance of 1089 Mpc, for $H_0$=75 km/s/Mpc and $q_0$ =0.5), this corresponds to 5pc. We expect a detectable LiH signal, since the H$_2$ column density is estimated to be N(H$_2$)$ = 5 \\times 10^{23}$ cm$^{-2}$. Menten \\& Reid (1996) derive an N(H$_2$) value ten times lower than this, using the H$_2$CO($2_{11}-2_{12}$) transition at 8.6\\,GHz. At this low frequency the structure and extent of the background continuum source may be quite larger than at 100--200 GHz and the source covering factor smaller. This means that their estimate of the column density is a lower limit.   \\section { Observations }  The observations were made with the IRAM 30m telescope at Pico Veleta near Granada, Spain. They were carried out in four observing runs, in December 1996, March, July and December 1997. Table 1 displays the observational parameters. We observed at 263 GHz with an SiS receiver tuned in single sideband (SSB). The SSB receiver temperature varied between 400 and 450K, the system temperature was 600-1400K depending on weather conditions, and the sideband rejection ratio was 10dB (the image frequency is at 271.5 GHz, in a region where the atmospheric opacity increases rapidly due to water vapour). We used a 512x1MHz filterbank and an autocorelator backend, with 0.3 km/s resolution. We present here only the 1MHz resolution spectra, smoothed to a 2.3 km/s channels, to improve the signal to noise.   \\begin{table} \\begin{flushleft} \\caption[]{ Parameters for the tentative LiH line } \\begin{tabular}{lccl} \\hline \\multicolumn{1}{l}{J$_u$--J$_l$ }                           & \\multicolumn{1}{c}{1--0 }              & \\multicolumn{1}{c}{}                   \\\\ $\\nu_{lab}$ GHz\t& 443.953 & \\\\ $\\nu_{obs}$ GHz\t& 263.527 & \\\\ Forward eff.\t  &  0.86 &\t\\\\ Beam eff.\t     &  0.32 & \\\\ T$_A^*$        & 7 mK\t& depth of absorption line \\\\ T$_{\\rm cont}$ & 15 mK & \\\\ FWHM      \t    & 3.2\tkm/s & \\\\ $\\sigma$       & 1.8 mK\t& noise rms with $\\Delta v$ 2.3 km/s \\\\ \\hline \\end{tabular} \\, \\\\ \\vskip 2truemm $\\alpha$(1950) = 02h 18m 04.1s  \\\\ $\\delta$(1950) = 35$^\\circ$ 42\\amin \\, 32\\asec  \\\\ \\end{flushleft} \\end{table}  The observations were done using a nutating subreflector with a 1' beamthrow in azimuth. We calibrated the temperature scale every 10 minutes by a chopper wheel on an ambient temperature load, and on liquid nitrogen. Pointing was checked on broadband continuum sources, and was accurate to 3\\asec \\, rms. The frequency tuning and sideband rejection ratios were checked by observing molecular lines towards Orion, DR21 and IRC+10216.  We integrated in total for 85 hours on the 263 GHz line, and obtained a noise rms level of 1.8 mK in the T$_A^*$ antenna temperature scale, with a velocity resolution of 2.3 km/s. The forward and beam efficiencies at the observed frequency are displayed in Table 1. The continuum level was estimated by observing in a rapid on--off mode using a special continuum backend. The switch frequency of the subreflector was increased from 0.5 Hz to 2 Hz.   \\begin{figure} \\psfig{figure=Ad212_f1.ps,bbllx=3cm,bblly=65mm,bburx=11cm,bbury=195mm,width=8cm} \\caption[]{ Spectrum of LiH in its fundamental line (1--0) at 444 GHz, redshifted at 263 GHz, in absorption towards B0218+357, compared to the highly optically thick CO(2--1) line previously detected. The tentative LiH line is slightly shifted from the center by about 5km/s, but is still comprised within the CO(2--1) velocity range. Its width is compatible with what is expected from an optically thin line. Spectra have been normalised to the absorbed continuum level and the velocity resolution is 2.3 km/s } \\label{lih_f1} \\end{figure} ",
        "conclusions": "Figure \\ref{lih_f1} presents our LiH spectrum, compared to that of CO(2--1) previously detected with the IRAM 30m-telescope (Wiklind \\& Combes 1995, Combes \\& Wiklind 1995). There is only a tentative detection of LiH at $\\sim$ 3 $\\sigma$. The line is very narrow, but is compatible to what is expected from an optically thin line. The CO(2--1) is highly optically thick, with $\\tau \\sim$ 1500. This optical depth is determined from the detection of C$^{18}$O(2--1), which is moderately thick, and the non--detection of C$^{17}$O(2--1). The center of the tentative line is shifted by 5 km/s from the average center of other lines detected towards B0218+357. This shift cannot be attributed to uncertainties of the line frequency, since it has been measured in the laboratory (e.g. Bellini et al 1994), and the error is at most 0.24 km/s at 3$\\sigma$, once redshifted. But the scatter of the line centers is $\\sim$ 3 km/s, and the width of most of the lines is $\\sim$ 15 km/s (cf Wiklind \\& Combes 1998). The velocity shift is therefore insufficient to reject the line as real.  Combining our own continuum data with that of lower frequencies (obtained from the NASA Extragalactic Database NED), we have previously found that the continuum spectra of B0218+357 can be fitted with a power law of slope --0.25 (Combes \\& Wiklind 1997). This would imply a continuum level of 15.5 mK at 263 GHz, which is in accord with the measured level. Since only 70\\% of the continuum is covered by molecular gas, the continuum level to be used for our LiH observations amounts to 11 mK.  \\smallskip  We can write the general formula, concerning the total column density of the LiH molecule, observed in absorption between the levels $l \\rightarrow u$ with an optical depth $\\tau$ at the center of the observed line of width $\\Delta v$ at half-power: $$ N_{LiH} = {{8\\pi}\\over{c^3}} f(T_x) {{\\nu^3 \\tau \\Delta v} \\over {g_u A_u} } $$ where $\\nu$ is the frequency of the transition, $g_u$ the statistical weight of the upper level ($= 2 J_u+1$), $A_u$ the Einstein coefficient of the transition, $T_x$ the excitation temperature, and $$ f(T_x) = {{Q(T_x) exp(E_l/kT_x)} \\over { 1 - exp(-h\\nu/kT_x)}} $$ where $Q(T_x)$ is the partition function. For the sake of simplicity, we adopt the hypothesis of restricted Thermodynamical Equilibrium conditions, i.e. that the excitation temperature is the same for all the LiH ladder. Since the line is not optically thick, but the optical thickness reaches $\\tau$ = 1.3 at the center of the line, we have derived directly from the spectrum, through a Gaussian fit of the opacity, the integrated $\\tau \\Delta v$  = 3.64 km/s. From the formulae above, and assuming an excitation temperature of $T_x$ = 15 K (see Table 2 for variation of this quantity), we derive a total LiH column density of 1.6 10$^{12}$ cm$^{-2}$ towards B0218+357. Compared to our previously derived H$_2$ column density of 5 10$^{23}$  cm$^{-2}$, this gives a relative abundance of LiH/H$_2$ $\\sim$ 3 10$^{-12}$. Note that there is a possible systematic uncertainty associated with this measure, due to the velocity difference between the maximum opacity of the CO, HCO$^+$ and other lines with that of LiH.  \\begin{table} \\begin{flushleft} \\caption[]{Derived LiH column density } \\begin{tabular}{lccccc} \\hline & & & & & \\\\ \\multicolumn{1}{c}{$T_x$ }                     & \\multicolumn{1}{c}{(K)}                   & \\multicolumn{1}{c}{5 }                   & \\multicolumn{1}{c}{10 }                   & \\multicolumn{1}{c}{15 }                   & \\multicolumn{1}{c}{20 }                  \\\\ & & & & & \\\\ \\hline & & & & & \\\\ N(LiH) & (10$^{12}$ cm$^{-2}$)  & 0.4 & 0.9 & 1.6 & 2.4 \\\\ & & & & & \\\\ LiH/H$_2$ & (10$^{-12}$) & 0.8  & 1.8 & 3.2 &  5 \\\\ & & & & & \\\\ \\hline \\end{tabular} \\end{flushleft} \\end{table}  \\smallskip  To interpret this result, comparison should be made with the atomic species. First, it is likely that the molecular cloud on the line of sight is dense and dark, and all the hydrogen is molecular, f(H$_2$) = 0.5. The Li abundance (main isotope $^7$Li) at $z=0.68466$ (i.e 5-10 Gyr ago) can be estimated at Li/H  $\\sim 10^{-9}$, since its abundance in the ISM increases with time. The primordial Li abundance must be similar to that in metal deficient unevolved Population II stars, Li/H = 1-2 10$^{-10}$ (Spite \\& Spite 1982), but Li could be depleted at the stellar surface by internal mixing. In meteorites and unevolved, unmixed Pop I stars, Li/H $\\sim 10^{-9}$, representative of the Li abundance some 4 Gyr ago. The present abundance in the ISM is estimated around 3 10$^{-9}$ (Lemoine et al 1993).  We therefore deduce LiH/Li $\\sim$ 1.5 10$^{-3}$. The uncertainty associated with the derived abundances are large, but the low LiH/Li ratio seems to exclude complete transformation of Li into LiH, as would be expected in very dense clouds (e.g. Stancil et al 1996, although the Li chemistry is not yet completely understood in dark clouds). However, it is likely that the cloud is clumpy, and in some of the more diffuse parts, LiH is photodissociated (e.g. Kirby \\& Dalgarno 1978). Also, some regions of the cloud could have a higher excitation temperature, in which case our computation under-estimates the LiH abundance (although the absorption technique selects preferentially cold gas, and the black-body temperature at the redshift of the absorbing molecules is $T_{bg}$ = 4.6 K).  The present observations suggest that the detection of LiH in emission towards dense clouds in the Milky Way should be easy with a submillimeter satellite, provided that the spatial resolution is enough to avoid dilution of the dense clumps. It is also interesting to observe the rarer molecule $^6$LiH, which in some clouds might be of same order of abundance as the main isotopic species. Through optical absorption lines Lemoine et al (1995) find towards two velocity components in $\\zeta$-Oph, $^7$Li/$^6$Li = 8.6 and 1.4.  Since $^6$Li is formed only in negligible amounts in the Big Bang, this ratio indicates that cosmic ray spallation has increased significantly the Li abundances.   \\vspace{0.25cm}"
    },
    "9804/astro-ph9804186_arXiv.txt": {
        "abstract": "We analysed 13 archival \\R\\ PSPC and HRI observations which included the position of a newly discovered 59\\,s X--ray pulsar in the Small Magellanic Cloud, 1SAX J0054.9--7226 = \\src. The source was detected three times between 1991 and 1996 at a luminosity level ranging from $\\sim$8$\\times$10$^{34}$ - 4$\\times$ 10$^{35}$ erg s$^{-1}$ (0.1--2.4 keV). Highly significant pulsations at 59.072\\,s were detected during the 1991 October 8--9 observation. The \\R\\ period, together with those measured by \\RXTE\\ and \\BSAX\\ yields a period derivative of \\.P= -- 0.016 s yr$^{-1}$. A much more accurate source position (10$^{\\prime\\prime}$ uncertainty) was obtained through the \\R\\ HRI detection on 1996 April restricting to three m$_ V$ $>$ 15.5 stars the likely counterpart of 1SAX J0054.9--7226 = \\src. ",
        "introduction": "On 1998 January 20 during a \\RXTE\\ observation in the direction of the Small Magellanic Cloud (SMC), a previously unknown X--ray source,  namely XTE J0055--724, was detected at a flux level (2--10 keV) of $\\sim$ 6.0 $\\times$ 10$^{-11}$ erg s$^{-1}$ cm$^{-2}$. The source showed pulsations at a period of $\\sim$59\\,s (Marshall \\& Lochner, 1998a). A previous \\RXTE\\ observation of the same field performed on 1998 January 12 failed to detect the source.  In response to these findings, simultaneous \\BSAX\\ and \\RXTE\\ observations of a region including the \\RXTE\\ error circle ($\\sim$10$^{\\prime}$ radius) of XTE J0055--724, were carried out on 1998 Jannuary 28. The results of these observations are reported elsewhere (Santangelo \\etal 1998a; Marshall \\etal 1998b). Thanks to the spatial capabilities of the imaging X--ray concentrators on board \\BSAX, an improved position ($\\sim$40$^{''}$ radius) was obtained for the pulsating source, named 1SAX J0054.9--7226 (Santangelo \\etal 1998a,b). The new \\BSAX\\ error circle contains only the previously classified \\R\\ and \\E\\ X--ray sources 1WGA J0054.9--7226 and \\src, which are likely the same source. In the following we adopt the earliest source name, i.e. \\E's. \\src\\ is a variable X--ray source in the SMC, which was already considered a candidate High Mass X--ray Binary by Wang \\& Wu (1992; source \\#35), Bruhweiler et al.  (1987; source \\# 9) and by White \\etal (1994; in the WGACAT), based on its high spectral hardness.  We report in this letter on the results from the analysis of the Position Sensitive Proportional Counter (PSPC) and High Resolution Imager (HRI) observations from the \\R\\ public archive. ",
        "conclusions": "\\src\\ was detected three times between 1991 and 1996 in the \\R\\ archival data. Highly significant pulsations, at a period of 59.072\\,s were detected on 1991 October 8--9. These findings, together with the \\BSAX\\ results, yield a mean period derivative of $\\sim$ -- 0.016\\, s yr$^{-1}$ between 1991 and 1998.  \\begin{figure}[tbh] \\centerline{\\psfig{figure=59s_dss.ps,width=8.cm,height=8.cm}} \\caption{ESO plate including the position of \\src. The X--ray error circles obtained from different instruments and satellites are shown} \\end{figure}  In one case a spectral analysis could be performed. The spectrum was found to be consistent with a relatively flat low absorbed power--law model that is typical of accreting X--ray pulsars in this energy range.  The 0.1--2.4 keV luminosity of \\src\\ as observed with \\R\\ ranges between $\\sim$4.2$\\times$10$^{35}$ erg s$^{-1}$ (1991 October 8--9) and $\\sim$8.5$\\times$10$^{34}$ erg s$^{-1}$ (1996 April 26 -- June 10). Moreover \\RXTE\\ detected \\src\\ at a luminosity level of $\\sim$3$\\times$10$^{37}$ erg s$^{-1}$ in the 2--10 keV energy band. Extrapolating to the \\R\\ energy range the luminosity measured by \\RXTE\\ on 1998 January 20, a 0.1--2.4 keV luminosity of $\\sim$2.5$\\times$10$^{36}$ erg s$^{-1}$ is derived, implying a pronounced long--term variability of \\src\\ (a factor of $>$30). This indicates that the source is probably a transient X--ray pulsar in a  high--mass binary containing a Be star.  A 10$^{\\prime\\prime}$ accurate position was obtained thanks to a \\R\\ HRI observation during which the source was detected (1996 April; 0.1--2.4 keV luminosity of $\\sim$8.5$\\times$10$^{34}$ erg s$^{-1}$). The \\R\\ HRI error circle of contains only three stars in the ESO plates with m$_V$ $>$ 15.5, the likely optical counterpart of \\src (see Fig.\\,3). Assuming a B--V = --0.2 and a distance modulus of 19 mag, these optical counterpart candidates are consistent with main sequence A9 -- B2 stars. We note that a similar spectral--type star (B1.5Ve; m$_V$ = 16) is the companion of the nearby X--ray source SMC X--2. Future optical follow--up observations of these candidates should determine the counterpart of \\src\\ and its probable Be star X--ray transient nature. The optical and/or infrared activity brightening of the counterpart will allow further X--ray triggers and studies."
    },
    "9804/astro-ph9804009_arXiv.txt": {
        "abstract": "The interstellar cloud surrounding the solar system regulates the galactic environment of the Sun, and determines the boundary conditions of the heliosphere. Both the Sun and interstellar clouds move through space, so these boundary conditions change with time. Data and theoretical models now support densities in the cloud surrounding the solar system of  n(H$^{\\circ}$)=0.22$\\pm$0.06 cm$^{-3}$, and n(e$^{-}$)$\\sim$0.1 cm$^{-3}$, with larger values allowed for n(H$^{\\circ}$) by radiative transfer considerations. Ulysses and Extreme Ultraviolet Explorer satellite He$^{\\circ}$ data yield a cloud temperature of {\\mbox 6,400 K}. Nearby interstellar gas appears to be structured and inhomogeneous. The interstellar gas in the Local Fluff cloud complex exhibits elemental abundance patterns in which refractory elements are enhanced over the depleted abundances found in cold disk gas. Within a few parsecs of the Sun, inconclusive evidence for factors of 2--5 variation in Mg$^{+}$ and Fe$^{+}$ gas phase abundances is found, providing evidence for variable grain destruction. In principle, photoionization calculations for the surrounding cloud can be compared with elemental abundances found in the pickup ion and anomalous cosmic ray populations to model cloud properties, including ionization, reference abundances, and radiation field. Observations of the hydrogen pile-up at the nose of the heliosphere are consistent with a barely subsonic motion of the heliosphere with respect to the surrounding interstellar cloud. Uncertainties on the velocity vector of the cloud that surrounds the solar system indicate that it is uncertain as to whether the Sun and $\\alpha$ Cen are or are not immersed in the same interstellar cloud. ",
        "introduction": "The physical conditions of the surrounding interstellar cloud establish the boundary conditions of the solar system and heliosphere. The abundances and ionization states of elements in the surrounding interstellar cloud determine the properties of the parent population of the anomalous cosmic ray and pickup ion components.  In addition, the history of the interstellar environment of the heliosphere appears to be partially recorded by radionucleotides such as $^{10}$Be and $^{14}$C in geologic ice core records (\\cite{sonett,fr97}). Because the solar wind density decreases as $R^{-2}$ ($R$=distance to Sun), the solar wind and interstellar densities are equal at about 5 AU (the orbit of Jupiter), in the absence of substantial ``filtration'' \\footnote{``Filtration'' refers to the deflection of interstellar H$^{\\circ}$ around the heliopause due to the coupling between interstellar protons and H$^{\\circ}$ resulting from charge exchange}. Approximately 98\\% of the diffuse material in the heliosphere is interstellar gas (\\cite{gruntman}). Thus, the physical properties of the outer heliosphere are dominated by interstellar matter (ISM).  Were the Sun to encounter a high density interstellar cloud, it is anticipated that the physical properties of the inner heliosphere would also be ISM-dominated. Zank and Frisch (1998) have shown that if the space density of the interstellar cloud which surrounds the solar system were increased to $\\sim$10 cm$^{-3}$, the properties of the inner heliosphere at the 1 AU position of the Earth would be dramatically altered.  The accuracy with which the physical properties of the surrounding cloud can be derived from observations of stars within a few parsecs of the Sun (1 pc$\\sim$200,000 AU) depends on the homogeneity and physical parameters of nearby ISM. Observations of nearby stars gives sightlines which probe the ensemble of nearby clouds constituting the ``Local Fluff'' cloud complex. Conclusions based on observations of nearby stars,  however, must be qualified by the absence of detailed data pertaining to the small scale structure of the local ISM (LISM).  More distant cold diffuse interstellar gas is highly structured, replete with dense ($\\sim 10^{4}-10^{5}$ cm$^{-3}$), small (20--200 AU) inclusions occupying in some cases less than 1\\% of the cloud volume (\\cite{frail,falgarone,falpug,heiles}). Small scale structures are ubiquitous in interstellar gas, and individual velocity components exhibiting column densities as low as N(H$^{\\circ}$)$\\sim$3$\\times 10 ^{18}$ cm$^{-3}$ are found in cold clouds (\\cite{frail,heiles}). The presence of dense low column density wisps near the Sun is allowed by currently available data.  The Sun has a peculiar motion with respect to the ``Local Standard of Rest'' (LSR\\footnote{The LSR is the velocity frame of reference in which the vector motions of a group of nearby comparison stars are minimized. Stars in the LSR corotate around the galactic center with a velocity of $\\sim$250 km s$^{-1}$}); the Sun moves through the LSR with a velocity V$\\sim$16.5 km s$^{-1}$ towards the apex direction l=53$^{\\circ}$, b=+25$^{\\circ}$ (\\cite{mihalas}). Uncertainties on the relative solar-LSR motion appear to be less than 3 km s$^{-1}$ and $\\pm$5$^{\\circ}$.  This motion corresponds to $\\sim$17 pc per million years.  Note that the solar path is tilted by $\\sim25^{\\circ}$ with respect to the galactic plane.  The Sun oscillates about the galactic plane, crossing the plane every 33 Myrs, reaching a maximum distance from the plane of $\\sim$77 pc.  The last galactic plane ``crossing'' was about 21 Myrs ago (\\cite{bash}). This amplitude of oscillation can be compared to scale heights on the order of $\\sim$50-80 pc for cold H$_{2}$ and CO, $\\sim$100 pc for cold H$^{\\circ}$ and infrared cirrus, $\\sim$250 pc for warm H$^{\\circ}$, and $\\sim$1 kpc for warm H$^{+}$ (the ``Reynolds Layer'').  There are three time scales of interest in understanding the environmental history of the solar galactic milieu -- $\\sim 10^{6}$ years, $\\sim 10^{5}$ years, and $\\sim 10^{4}$ years. Prior to entering into the Local Fluff complex of interstellar clouds, the Sun traveled through a region of the galaxy between the Orion spiral arm and the spiral arm spur known as the Local Arm.  On the order of a million years ago, the Sun was displaced $\\sim$17 pc in the anti-apex direction, which is towards the present day location of the junction of the borders of the constellations of Columba, Lepus and Canis Major. The motions of the Sun and surrounding interstellar cloud with respect to interstellar matter within 500 pc, projected onto the plane, are illustrated in Figure 1.  Note that the velocity vectors of the Sun and interstellar cloud surrounding the solar system are nearly perpendicular in the LSR, implying that the surrounding cloud complex is sweeping past the Sun (see section \\ref{velocity}). When the morphology of the Local Fluff complex is considered, it is apparent that sometime during the past $\\sim$200,000 years the Sun appears to have emerged from a region of space with virtually no interstellar matter (densities n(H$^{\\circ})<0.0005$ cm$^{-3}$, n(e$^{-})<0.02$ cm$^{-3}$) and entered the Local Fluff complex of clouds (average densities n(H$^{\\circ}$)$\\sim$0.1 cm$^{-3}$) outflowing from the Scorpius-Centaurus Association of star-forming regions. One model for the morphology of the cloud surrounding the solar system predicts that sometime within the past 10,000 years, and possibly within the past 2,000 years, the Sun appears to have entered the interstellar cloud in which it is currently situated (\\cite{fr94}, Frisch 1997). The cloud surrounding the solar system will be called here the ``Local Interstellar Cloud'' (LIC\\footnote{This cloud surrounding the solar system is also referred to as the ``surrounding interstellar cloud'', or SIC, which unambiguously defines the cloud feeding interstellar matter into the solar system.  For the sake of uniformity of notation, however, the term LIC is used here.}).  \\begin{figure} \\begin{center} \\plotone{fig1.eps} \\end{center} \\vspace{0.5in} \\caption[]{{\\small The distribution of interstellar molecular clouds (traced by the CO 1-$>$0 115 GHz rotational transition) and diffuse gas (traced by E(B-V) color excess due to the reddening of starlight by interstellar dust) within 500 pc of the Sun are shown.  The round circles are molecular clouds, and the shaded material is diffuse gas. The horizontal bar (lower left) illustrates a distance of 100 pc. Interstellar matter is shown projected onto the galactic plane, and the plot is labeled with galactic longitudes.  The distribution of nearby interstellar matter is associated with the local galactic feature known as ``Gould's Belt'', which is tilted by about 15--20$^{\\circ}$ with respect to the galactic plane. ISM towards Orion is over 15$^{\\circ}$ below the plane, while the Scorpius-Centaurus material (longitudes 300$^{\\circ}$--0$^{\\circ}$) is about 15--20$^{\\circ}$ above the plane.  Also illustrated are the space motions of the Sun and local interstellar gas, which are nearly perpendicular in the LSR velocity frame. The three asterisks are three subgroups of the Scorpius-Centaurus Association.  The three-sided star is the Geminga Pulsar.  The arc towards Orion represents the Orion's Cloak supernova remnant shell.  The other arcs are illustrative of superbubble shells from star formation in the Scorpius-Centaurus Association subgroups.  The smallest (i. e. greatest curvature) shell feature represents the Loop I supernova remnant.}} \\label{fig1} \\end{figure} ",
        "conclusions": "One new conclusion presented here is that in principle the {\\it in situ} pickup ion data can help resolve the outstanding question of whether the correct reference abundances for the LIC are given by solar versus B-star abundances.  A second new result is that the uncertainties on the LIC velocity vector indicate that it is not yet clear whether the Sun and $\\alpha$ Cen are immersed in the same interstellar cloud.  Based on the discussions in this paper, the best values for LIC properties are given by n(H$^{\\circ}$)=0.22$\\pm$0.06 cm$^{-3}$, n(e$^{-}$)=n(H$^{+}$)=0.1 cm$^{-3}$, T=6,900 K and a relative Sun-cloud velocity of 25.8$\\pm$0.8 km s$^{-1}$. However, radiative transfer considerations in the LIC suggest that the quoted neutral density is a lower limit. Ulysses and EUVE observations of He$^{\\circ}$ indicate a cloud temperature of T=6,400 K. The magnetic field strength is weakly constrained to be in the range of 2--3 $\\mu$G.  Models of the Ly$\\alpha$ absorption line towards $\\alpha$ Cen are consistent with an Alfven velocity of 20.9 km s$^{-1}$, which in turn is consistent with an interstellar magnetic field of 3 $\\mu$G in the absence of additional unknown contributions to the interstellar pressure. Ulysses and EUVE observations of interstellar He$^{\\circ}$ within the solar system give an upwind direction for the ``wind'' of interstellar gas through the solar system, in the rest frame of the Sun, of V=--25.9$\\pm$0.6 km s$^{-1}$ arriving from the galactic direction l=4.0$^{\\circ}$$\\pm$0.2$^{\\circ}$, b=15.4$^{\\circ}$$\\pm$0.6$^{\\circ}$. Removing solar motion from this vector gives an upwind direction for the LIC cloud in the LSR of V=--18.7$\\pm0.6$ km s$^{-1}$ arriving from the direction l=327.3$^{\\circ}$$\\pm$1.4$^{\\circ}$, b=0.3$^{\\circ}$$\\pm$1.0$^{\\circ}$.  Through a combination of observations and theory, uncertainties in the LIC electron density are narrowing. Radiative transfer in the sightlines towards nearby stars require that cloud models must be combined with data in order to deduce properties at the cloud location. Radiative transfer models of ionization in the LISM show interesting results, but additional understanding of the input radiation fields is needed. The Local Fluff complex is structured and inhomogeneous. Striking progress would be made in understanding this structure if interstellar absorption lines could be observed at resolutions of $\\sim$1 km s$^{-1}$ in the ultraviolet. The most glaring uncertainty is the absence of detailed knowledge about the interstellar magnetic field. Many of the most abundant elements in the LIC are ionized, and densities of neutral atoms with FIPs less than 13.6 eV are typically down by 1--3 orders of magnitude from the dominant ions. The current approach of trying to understand the interaction of the ISM with the heliopause, from both the outside in and the inside out, is finally bearing fruit."
    },
    "9804/astro-ph9804192_arXiv.txt": {
        "abstract": "The nuclei of a wide class of active galaxies emit broad emission lines with widths at half maximum (FWHM) in the range $10^{3}-10^{4}$ km s$^{-1}$.  This spread of widths is not solely a consequence of the range of the luminosities of these sources since a plot of width versus luminosity shows a large scatter. We propose that the broad line emission region (BLR) is axially symmetric and that this scatter in line width arises from an additional dependence on the angle of the line of sight to the axis of the emission region.  Such a relation is natural in unified models of active nuclei which link a variety of observed properties to viewing angle. Adopting a simple form for the line width as a function of luminosity and angle, and convolving this with the observed luminosity function, allows us to predict a line width distribution consistent with the available data. Furthermore, we use the relation between the equivalent width of a line and the luminosity in the continuum (the `Baldwin Effect') to predict an observed correlation between line width and equivalent width. The scatter on this correlation is again provided by angular dependence. The results have applications as diagnostics of models of the broad line emission region and in cosmology. ",
        "introduction": "In unified models of active galactic nuclei with spherically symmetric BLR the width distribution of the broad emission lines cannot be accounted for by luminosity dependence alone. Plots of line width versus continuum luminosity have a large scatter and show no significant correlation \\cite{W93,P97}. There is, however, growing evidence that the broad line region (BLR) is not spherical, but axisymmetric. \\begin{enumerate} \\item Observed samples of AGN \\cite{Wills86,W93,Brotherton94} suggest relations between line widths and R, the ratio of core to lobe dominance. Other samples \\cite{P97} find relations between line widths and $\\alpha_{\\rm ox}$, the continuum slope parameter from the optical to X--ray bands. Both of these parameters have some viewing angle dependence. \\item The continuum and line light curves of some active nuclei, eg 3C390.3 \\cite{Wamsteker97}, are most naturally interpreted in terms of a disc-like line emission region. It has also been suggested that some double peaked line profiles arise from discs (eg Arp 102b) although the interpretation in these cases is not so clear when time variability is taken into account. \\item Several axisymmetric disc-wind models, such as those of Cassidy \\& Raine \\shortcite{Cassidy96}, Chaing \\& Murray \\shortcite{Chaing96} and Emmering, Blandford \\& Shlosman \\shortcite{Emmering92} have been proposed and  models of this type are gaining support from evidence for winds \\cite{Pasadena}.  These will naturally predict some viewing angle dependence of line width. \\end{enumerate} It should be noted that Osterbrock \\shortcite{Osterbrock77} showed that a deficit of systems with narrower lines ruled out pure disc models, but such objections do not necessarily apply to axisymmetric models in general.  In this paper we shall adopt a simple dependence of line width on both viewing angle and luminosity. Then: \\begin{enumerate} \\item We obtain a reasonable fit to the distribution of line widths. \\item Given the Balwin relation between line and continuum luminosity we predict a relation between line width and equivalent width compatible with the observed trend. The scatter on this relation is attributed to angular dependence. \\item We discuss how the width distribution can be used to test models of the BLR. \\item If the BLR is indeed axisymmetric, we show how the line width distribution can be used, in principle, to determine cosmological parameters. \\end{enumerate} ",
        "conclusions": "We conclude that the simple picture we have presented here accounts for the  scatter in FWHM versus luminosity, accounts for the distribution of FWHM, and relates the trend in the EW with FWHM to the Baldwin relation. This may be useful as a diagnostic tool in discriminating between disc-wind models. The analysis has applications as a cosmological tool particularly as measurement of line widths is independent of any cosmological model."
    },
    "9804/astro-ph9804317_arXiv.txt": {
        "abstract": "Collimated outflows from Young Stellar Objects (YSOs) can be seen as tracers of the accretion powered systems which drive them.  In this paper I review some theoretical and observational aspects of YSO outflows through the prism of questions relating to the protostellar source.  The issue I address is: can collimated outflows be used as ``fossils'' allowing the history of protostellar evolution to be recovered?  Answering this question relies on accurately identifying where theoretical tools and observational diagnostics converge to provide unique solutions of the protostellar physics.  I discuss potential links between outflow and source including the time and direction variability of jets, the jet/molecular outflow connection, and the the effect of magnetic fields.  I also discuss models of the jet/outflow collimation mechanism. ",
        "introduction": " ",
        "conclusions": "The issues cited in this paper are associated with outflows. How do these issues specifically relate to questions inherent to the physics of accretion? The time variability of jets relates to the time-dependence of accreation, the FU Ori outbursts being a notable example.  The direction variability of jets relates to the global dynamics and stability of accretion disks.  Livio \\& Pringle 1997, for example, have shown that radiation induced warping of disks may lead to precession in magneto-centrifugal jets. The presence and structure of magnetic forces in jets relates to the existence and form of large scale fields in the disks. If nose-cones do not occur in real YSO jets then perhaps mechanisms which rely on strong toroidal fields are excluded.  Thus YSO jets and outflows offer a unique opportunity for the study of accretion powered systems.  Protostellar outflows can be observed with exquiste detail in a variety of wavelengths including diagnostic spectral lines.  The quality of the data combined with the long lookback time inherent to the outflows offers the possibilty that a large fraction of individual protostar's history might be recovered if we learn were and how to look.  We are a long way from this now but the prospect of having such capabilities is very exciting."
    },
    "9804/astro-ph9804121_arXiv.txt": {
        "abstract": " ",
        "introduction": "A long way has been run from the first views developped to explain the formation of galaxies. In 1962, Eggen, Lynden-Bell \\& Sandage designed the collapse scenario, where all galaxies are created with their morphological type, according to their angular momentum. Their potentials remained axisymmetric, so that no angular momentum could be redistributed through gravity torques; the total mass and gas content was already there at first collapse. For elliptical galaxies, the violent/single collapse picture still remains in some modified form, although the most developped and adopted scenario is through agglomeration of a large number of clumps (e.g. van Albada 1982, Aguilar \\& Merritt 1990), that produces de Vaucouleurs profiles in $r^{1/4}$. The merger picture (Toomre 1977, Schweizer 1990), where ellipticals are formed by progressive interaction and coalescence of many parent galaxies, is favored in hierarchical cosmogonies.  For spiral galaxies, the scenario involves now much more internal dynamical evolution. Due to gas dissipation and cooling, gravitational instabilities are continuously maintained in spiral disks, and they drive evolution in much less than a Hubble time. Spiral galaxies are open systems, that accrete mass regularly, and their morphological type evolves along the Hubble sequence. Non-axisymmetric perturbations, such as bars or spirals, produce gravity torques that drive efficient radial mass flows; vertical resonances thicken disks and form bulges, and the mass central concentration can destroy bars. Accretion of small companions can also disperse bars and enlarge the bulge. A major merger can destroy disks entirely and form an elliptical.  The first role of galaxy interactions is to trigger internal evolution, that we will consider now, in the next section. Specific aspects of galaxy interactions and mergers will then be detailed in section \\ref{envir}. ",
        "conclusions": ""
    },
    "9804/astro-ph9804298_arXiv.txt": {
        "abstract": "We have produced radio maps, using the ATCA, of the central regions of six southern Seyfert 2 galaxies (NGC 1365, 4945, 6221, 6810, 7582, and Circinus) with circumnuclear star formation, to estimate the relative contribution of star formation activity compared to activity from the active galactic nucleus (AGN). The radio morphologies range from extended diffuse structures to compact nuclear emission, with no evidence, even in the relatively compact sources, for synchrotron self--absorption. In each case the radio to far--infrared (FIR) ratio has a value consistent with star formation, and in all but one case the radio to [FeII] ratio is also consistent with star formation. We derive supernova rates and conclude that, despite the presence of a Seyfert nucleus in these galaxies, the radio, FIR, and [FeII] line emission are dominated by processes associated with the circumnuclear star formation (i.e. supernova remnants and H~II regions) rather than with the AGN. ",
        "introduction": "Circumnuclear star formation is common in Seyfert galaxies, but the relationship between the Seyfert nucleus and the surrounding star formation is not well understood (see Filippenko 1992), and both evolutionary and causal relationships have been suggested. For example, a nuclear starburst may develop into a massive star cluster or black hole (Norman \\& Scoville 1988), or conversely the energy output from an active galactic nucleus (AGN) may trigger circumnuclear star formation (Sanders \\& Bania 1976). More recently, studies have indicated that star formation is occurring in and around the torus (e.g. Cid Fernandes \\& Terlevich 1992; Davies \\etal 1997) that is thought to dictate the type of Seyfert nucleus we observe.  Dust is a common feature of the circumnuclear environment of active galaxies. At optical wavelengths, it can obscure our view of the nucleus and hide any evidence of an AGN. However, radio observations are not only unaffected by dust but also have the advantage of high spatial resolution, and can be important in understanding the role of the various processes in active galaxies. For example, starburst galaxies generally have diffuse radio morphologies dominated by synchrotron emission from cosmic rays accelerated by supernovae, while Seyfert galaxies sometimes exhibit well--collimated radio jets and a compact nuclear radio source (see review by Condon 1992).  High-resolution radio synthesis images, with sub-arcsec resolution,  have found compact sources at the nucleus, and extended emission consisting of radio jets, ouflows, diffuse emission, and discrete sources such as HII regions and SNRs (e.g. Ulvestad \\& Wilson 1984; Kronberg \\etal 1985; Antonucci \\& Ulvestad 1988; Carral, Turner \\& Ho 1990; Condon \\etal 1991; Forbes \\etal 1994; Collison \\etal 1994; Sandqvist \\etal 1995).  High spatial resolution, multi--frequency radio observations can also be used to test the idea of advection-dominated accretion flows around black holes. For such flows the radio emission depends strongly on the mass of the black hole and is characterised by an inverted radio spectrum (Yi \\& Boughn 1997).  Here we present 3 and 6cm radio continuum images from the Australia Telescope Compact Array (ATCA) of 6 such galaxies which show evidence for narrow high excitation optical forbidden lines, and are classified as Seyfert 2s. In the case of NGC 1365, broad lines have also been detected from the nuclear region. All six galaxies show evidence for circumnuclear star formation, and several are barred. We discuss the radio morphologies of these galaxies and possible emission mechanisms, and we compare the radio data with that from other wavelengths to assess the relative contributions to the radio flux from star formation and the AGN. ",
        "conclusions": "In Figures 1 to 12 we show the 3 and 6cm radio images for the six galaxies in our sample. These galaxies do not represent a complete sample in any sense, but rather were chosen as relatively well--known southern Seyfert 2 galaxies that generally lacked high resolution radio maps. Our radio data for the Circinus galaxy have been presented elsewhere along with near--infrared line imaging (Davies \\etal 1997), although we include it here for comparison purposes with the other galaxies.  The images show a variety of radio morphologies which range from Circinus, with a strong, compact nucleus, to NGC 1365, with a extended region of discrete sources or hot-spots. The beam shape is shown in the lower left of each image.  Care is needed when interpreting these images, as (a) our observations are optimised for studying the nuclear region, and so our images will not correctly represent the extended diffuse emission from the disk of the galaxy, and (b) most of the galaxies have high inclinations so that confusion effects may play some role in the observed radio morphology.  Flux measurements in 2 arcsec and 6 arcsec diameter apertures for each galaxy are given in Table 2. We also list the 6cm to 3cm spectral index after smoothing the 3cm image to match the 6cm resolution. The spectral indices (F $\\propto \\nu ^{\\alpha}$) range from flat ($\\alpha$ $\\sim$ 0) to steep ($\\alpha$ $\\sim$ --1). The total 6cm flux in Table 2 is given both for our images and for the single--dish observations by Wright \\etal (1994, 1996). The difference between these indicates the amount of  diffuse emission which is missing from our images. We also calculated the maximum brightness temperatures for each image (both 3 and 6 cm) and the maximum value is given for each source.  In Table 3 we list various derived quantities for the sample including Hubble type, distance, inclination, 6cm and [FeII] line luminosities and the SN rate. The [FeII] measurements are from Moorwood \\& Oliva (1988) in a 6 arcsec diameter aperture, except for the Circinus galaxy in which we use a nuclear [FeII] flux from Davies \\etal (1997). None of the [FeII] fluxes have been corrected for extinction. The SN rate is calculated separately from both the 6cm and [FeII] line flux in the 6 arcsec aperture (corresponding to $\\sim$ 500 pc at a typical distance of 20 Mpc). The 6cm SN rate is calculated assuming that all of the 6cm flux is non--thermal emission from cosmic rays accelerated by SNRs (e.g. Condon \\& Yin 1990). This may give an overestimate of the rate because it ignores any nuclear flux (which may be significant) and the contribution from thermal emission (which is unlikely to be significant). For the [FeII] SN rate we simply assume an average luminosity of 2 $\\times$ 10$^{36}$ erg s$^{-1}$ over an adiabatic lifetime of 2 $\\times$ 10$^{4}$ yrs (e.g. Norris \\& Forbes 1995). For most galaxies the two methods give rates within a factor of two, the notable exception being NGC 4945 (which is discussed further below).   \\subsection{Radio Spectral Indices}  The radio spectral index of Seyfert and starburst galaxies is determined by four mechanisms.  1. Cosmic rays within the galaxy are generated and re--accelerated by supernovae and supernova remnants, and then interact with the interstellar magnetic field to emit synchrotron emission. This synchrotron emission, with a typical spectral index of $\\alpha$ $\\sim$ --0.7, is expected to dominate the radio power of starburst galaxies, and should appear as a diffuse component in radio images of these sources.  2. Relativistic particles ejected from the massive black hole at the nucleus of a galaxy may generate  intense synchrotron emission, similar to that seen in radio--loud galaxies and quasars. The cores in these radio-loud objects typically have a flat--spectrum core, indicating synchrotron self--absorption, and steep--spectrum extended radio--lobes, suggesting cooling of high--energy electrons. However, synchrotron self--absorption is important only for brightness temperatures greater than 10$^{10}$ K (Condon 1992). Most Seyfert galaxies, on the other hand, are observed to have much lower brightness temperatures than this in the core, so that synchrotron self-absorption is not significant in these sources. This is confirmed by the observed core spectral indices, which are frequently in the region of --0.7.  3. When the radiative efficiency in the accretion disk is low, an advection-dominated accretion flow (ADAF) may operate. The radio emission in these sources is dominated by synchrotron emission from a hot plasma, and the emission from such flows is predicted to give rise to inverted spectra with typical indices of +0.4 (Yi \\& Boughn 1997). The ADAF radio emission mechanism has only recently been proposed and should be regarded as untested at this stage. The data here are unlikely to provide a definitive test because of insufficient resolution, and all radio spectral indices measured here are negative. We will therefore not consider this mechanism further, except to note that this mechanism would be indicated by inverted-spectrum emission from a low-brightness-temperature core. No source here falls into this category.  4. H II regions in our galaxy generate free--free emission from hot electrons. Most are optically thin, giving a flat spectrum, although some compact H II regions become optically thick at centimetre wavelengths, giving a spectral index $\\sim$+2. However, the integrated flux of such regions is generally insignificant compared to the synchrotron emission of the host galaxy.  5.  The radio emission from ultra--luminous infrared galaxies is optically thick to free--free absorption, so that the typical synchrotron spectrum of these galaxies is flattened at low frequencies (Condon et al. 1991).  The combined result of these effects in Seyfert and starburst galaxies is to produce a typical radio spectral index of --0.7 (from the extended synchrotron emission) with a flattening at low frequencies in some starburst sources because of free--free absorption.  Table 2 shows that the nuclear spectral indices of three of the galaxies (NGC 1365, NGC 6221, NGC 7582) is --0.5 or steeper on both the 2 arcsec and the 6 arcsec scale, showing evidence for neither free--free nor synchrotron absorption. In the other three galaxies, the cores have flatter spectra, but the brightness temperatures ($\\le$ 8300 K) are too low for synchrotron self--absorption, indicating that free--free absorption is responsible for the flattening.  Of course, we cannot rule out the presence of a weak synchrotron self--absorbed core in the nucleus of any of these galaxies. However, comparison of the radio fluxes in a 2--arcsec aperture with the flux in a 6--arcsec aperture in Table 2 shows that the luminosity of any such core is small compared to the surrounding emission. Therefore any such core does not contribute significantly to the overall energy budget of the nuclear region of the galaxy, and is not responsible for the overall flat spectrum of the nuclear region..  This degree of free-free absorption flattening indicates either a high star formation rate (Condon \\etal 1991) or that we are viewing the AGN through optically--thick narrow-line-region clouds (Roy \\etal 1994).  \\subsection{The Radio - [FeII] Correlation}  Forbes \\& Ward (1993) discovered that the 6cm radio emission in the central regions of active galaxies is strongly correlated with the near--infrared [FeII] 1.64$\\mu$m line emission. This relation exists over several orders of magnitude. With a larger sample, Simpson \\etal (1996) were able to show that Seyfert and starburst galaxies follow slightly different radio--[FeII] relations. For starburst galaxies the relation, with slope $\\sim$ 1, can be reasonably explained by SNRs which are responsible for both the non--thermal radio emission and the fast shocks that provide the [FeII] excitation. However, the situation for Seyfert galaxies (which reveal a correlation slope of $\\sim$ 0.7) is less clear. Simpson \\etal argued that photo-ionisation from the Seyfert nucleus can cause this relationship, with a contribution from radio--jet induced shocks in some cases.  In Fig. 13 we show the [FeII]/6cm ratio for our sample galaxies, compared with the Seyfert and starburst relations of Simpson \\etal (1996). The 1$\\sigma$ dispersion of galaxies about the relations is $\\sim$ $10^{0.5}$. For the Seyferts studied here, we find a large degree of star--formation activity compared to photo--ionisation from an AGN, and so we might expect them to lie closer to the starburst relation than the Seyfert one. This indeed appears to be the case for four galaxies, although one (NGC 7582) is closer to the Seyfert relation and NGC 4945 falls well away from either relation. The [FeII]/6cm ratio of NGC 4945 is about a factor of 100 lower than typical active galaxies, and we discuss this further in Section 3.4 below. We note however that given the dispersion in the relations, and the low luminosities ($\\le$ 10$^{40}$ erg s$^{-1}$) of the galaxies studied here, this is not a sensitive test of the excitation mechanism.  \\subsection {The Radio -- FIR Correlation}  Normal spiral and starburst galaxies show a tight correlation between their radio and FIR luminosity (e.g. Wunderlich et al. 1987). This correlation, which extends over five orders of magnitude, is true for both flux density and luminosity, and cannot be accounted for by selection effects, or by a simple ``richness effect''. While a detailed mechanism to explain this correlation has yet to be established, it is almost certainly the result of star formation, which generates both the synchrotron radio emission and the thermal FIR emission. This is supported by the fact that all objects that are dominated by star formation (HII galaxies, normal spirals, starburst galaxies) do follow the correlation.  On the other hand, Sopp \\& Alexander (1991) showed that radio--loud quasars and radio galaxies clearly do not follow the radio--FIR correlation. Thus whether or not a galaxy follows this correlation may be used as an indicator of the dominant radio luminosity source of the galaxy.   Norris \\etal (1988) and Roy \\etal (1997) showed that Seyfert galaxies, unlike radio--loud quasars, do roughly follow the radio--FIR correlation, but with a looser fit than normal spirals and starbursts. This suggests that the bolometric luminosity of Seyfert galaxies may be dominated by star formation. This is supported by off--nuclear optical and infrared observations of Seyferts, which show the same line ratios and luminosities as starburst galaxies (Bransford et al. 1997). Thus, although the nuclear optical spectra of Seyfert galaxies are clearly dominated by an AGN, in many cases the integrated radio emission and the FIR emission are dominated not by the AGN but by circumnuclear star formation.  The degree to which an individual galaxy follows this correlation is most conveniently expressed by the parameter q -- the logarithm of the FIR to radio ratio. The conventional definition of q follows that of Helou \\etal (1985), who define q in terms of the 1.49 GHz radio flux. For our purposes, we adapt Helou's definition to our observing frequency of 4.8 GHz by assuming a spectral index of --0.7, and therefore define it as  \\hspace{.5in}q $\\equiv$ log[(FIR/3.75 x 10$^{12}$ Hz)/(2.26 x S$_{\\rm 4.8 GHz}$]\\hspace{1.5in}(1)  \\hspace{.5in}where FIR $\\equiv$ 1.26 x 10$^{-14}$(2.58S$_{\\rm 60\\mu}$+S$_{\\rm 100\\mu}$)\\hspace{2in}(2)  Typical values of q from the IRAS Bright Galaxy Sample are 2.34 for normal spirals, 2.21 for starburst galaxies, and less than 2 for radio--loud AGNs (Condon \\etal 1991).  All the galaxies studied here except NGC 4945 have q in the range 2.2 to 2.3, which places them firmly in the middle of the radio--FIR correlation, and suggests that most of their radio and FIR luminosity is produced by star formation. We discuss the case of NGC 4945 (q = 1.88) below.   \\subsection{Individual Galaxies}  Here we discuss each galaxy in turn, starting with an extended discussion of NGC 4945. To avoid repetition, we note that in every case other than NGC 4945, the spectral index, radio--FIR ratio, and [FeII]--radio ratio are all consistent with star formation, rather than an AGN, being the dominant source of radio emission.  \\noindent {\\bf NGC 4945} This infrared luminous galaxy is nearly edge--on and is located in a nearby loose group. Although we list it as a barred galaxy in Table 3, there is a continuing debate about the reality of the bar (e.g. Harnett \\etal 1989). Koornneef (1993) described NGC 4945 as a post--starburst galaxy with no evidence for an AGN. However Moorwood \\& Oliva (1994) have argued that the central regions do show signs of ongoing young star formation. Evidence for a heavily obscured AGN now come from the variable hard X--rays (Iwasawa \\etal 1993), and the presence of a compact radio core in VLBI observations (Sadler \\etal 1995). The galaxy contains a thick torus or ring  with a radius of $\\sim$ 150 pc (Koornneef 1993; Moorwood \\etal 1996). Harnett \\etal (1989) found that the radio emission has a strong central contribution with  emission extended 10 arcmin perpendicular to the major axis. Multi--frequency observations have been carried out by Elmouttie \\etal (1997). They focused on the large scale structure using a beam size of $\\sim$ 25$^{''}$, and found that the spectral index steepens from the central region to the main disk of the galaxy.  Furthermore, NGC 4945 is notable for the fact that it is one of the few galaxies (along with Circinus) known to contain water megamasers. Such megamasers have been cited in NGC 4258 (Miyoshi \\etal  1995) as the strongest evidence known for a massive black hole in an AGN. Preliminary VLBI imaging (Greenhill \\etal 1997) of the megamasers supports the model that they are in a Keplerian disk surrounding the black hole.  Our radio image is dominated by strong nuclear emission, and emission extended along the disk of the galaxy. However there is also evidence of some filamentary structure perpendicular to the major axis. Such extended emission may be associated with the outflowing superwind in this galaxy (Nakai \\etal 1989; Lipari, Tsvetanov \\& Macchetto 1997). The extended emission has a steeper spectral index ($\\alpha \\sim -0.8$) than the radio nucleus, however we note that our data are less sensitive to extended emission (particularly at 3cm) which makes the spectral index somewhat less certain. The nucleus has a relatively flat spectral index of $\\alpha$ = --0.3. The brightness temperature in the central 2 arcsec of this source (i.e. 7000 K) is still far too low for synchrotron self--absorption, indicating that the star formation activity is particularly intense, to provide the necessary free--free absorption. The obscuration inferred from the X--ray data suggest that the extinction towards the AGN could be as high as A$_V$ $\\sim$ 2500.  We noted above (in Section 3.3) that the radio--FIR ratio for this galaxy is unusually low (i.e. q = 1.88), which at first sight appears to suggest that an AGN is responsible for much of the radio emission. However, this galaxy is so near that not all the FIR flux was in the single IRAS aperture, and so the IRAS flux listed in the IRAS Point Source Catalog may be an underestimate. Rice \\etal (1988) have estimated the total FIR flux by co--adding IRAS images and obtain a higher value for the FIR fluxes, which raises the value of q to 2.1, suggesting that the radio emission in this galaxy is again dominated by star formation rather than by an AGN.  An interesting property of NGC 4945 is that the ratio of the [FeII] line luminosity to 6cm radio emission is only 0.63, which is almost a factor of 100 less than is typical for active galaxies (see Fig. 13). We now consider a number of possible reasons for this.  \\begin{itemize}  \\item The reduced ratio could be due to extinction (by dust) of the [FeII]. However,  Moorwood \\& Oliva (1994) estimate that the extinction in the [FeII] line emitting zone is 1.8 mag or a factor of five, which is insufficient to produce the observed effect.  \\item It could be because of nuclear radio emission from an AGN which is not accompanied by [FeII] line emission. We have shown above that the radio/FIR ratio for the galaxy as a whole is consistent with star formation activity. However, the central 6 arcsec (over which we measure the [FeII]/6cm ratio) contributes only 9\\% of the total radio flux, and we have no information on the radio/FIR ratio in the nucleus, so the radio flux from the AGN could be abnormally large. In this case, we would expect the [FeII]/6cm ratio to approach the usual value as we increase the area over which we integrate the flux (for both 6 cm and [FeII]). However, Moorwood \\& Oliva (1994) quote a total [FeII] flux over an emitting region of 18 $\\times$ 21 arcsec to be 12 $\\times$ 10$^{-14}$ erg s$^{-1}$ cm$^{-2}$, or an observed log luminosity of 38.81 erg s$^{-1}$. The 6cm radio luminosity over a similar area is 40.21 erg s$^{-1}$, giving a [FeII]/6cm ratio of 0.03, which is even lower than the value in the nucleus, indicating that the [FeII]/6cm ratio falls off with distance from the galaxy centre, and that the low value is not a consequence of nuclear radio emission.  \\end {itemize}  Thus the abnormally low [FeII]/6cm ratio in NGC 4945 of 0.63 is produced in the region surrounding the nucleus, where the radio emission (with a spectral index of $\\sim$ --0.7) is due to SNRs in the galaxy disk, the outflowing starburst superwind discussed above, or perhaps a radio jet. Pure SNRs produce ratios of about 500 i.e. well in excess of typical galaxy values, so this would tend to give an enhanced ratio. In a $6^{''} \\times 6^{''}$ aperture, the superwind galaxies M82 and NGC 253 have [FeII]/6cm ratios of 75 and 52 respectively, although the superwind itself in NGC 253 does not seem to produce significant [FeII] line emission (Forbes \\etal 1993). Again such ratios are significantly higher than seen in NGC 4945. The data for Seyfert galaxies with clear radio jets are limited. For NGC 4151 and NGC 1068 the measured ratios are 28 and 9. This is closer to the NGC 4945 value but still a factor of at least 10 too high. We conclude that the abnormal [FeII]/6cm ratio in NGC 4945 is due to either (a) a starburst superwind, which produces substantial radio emission but very little [FeII] line flux (due perhaps to an unknown excitation effect or low density in the wind), or (b) a radio jet which dominates the extended radio emission on the few--arcsec scale but which does not produce significant [FeII] emission.  \\noindent {\\bf NGC 1365} This is a well--studied barred galaxy. The central region reveals broad and narrow emission lines (Veron \\etal 1980) surrounded by a circumnuclear ring of star formation (Edmunds \\& Pagel 1982; Saikia \\etal 1994). The star formation, combined with the obscuring effects of dust, give the appearance of a hot-spot nucleus (Sersic \\& Pastoriza 1965). A high excitation outflow from the nucleus has been seen (e.g. Hjelm \\& Lindblad 1996).  High resolution radio continuum observations have been reported by several workers (e.g. Sandqvist, Jorsater \\& Lindblad 1982, 1995). In particular,  Sandqvist, Jorsater \\& Lindblad (1995) observed it with the VLA at  20, 6, and 2 cm. Their radio images revealed a weak nucleus surrounded by a elongated $\\sim$ 8 $\\times$ 20 arcsec (a/b = 0.4) ring of hot-spots, or components. They labelled a number of components A--H, of which B is blended with A and C is blended with D at $\\ge$ 1 arcsec resolutions. Our radio image, shown in Fig. 1, is consistent with theirs, except that we identify one additional component to the SW, which we call `J'. The hot-spots generally have steep spectra with 6cm luminosities of $\\sim$ 10$^{36}$ erg s$^{-1}$ which suggests that the radio emission from each component is made up of several SNRs. A combined radio and X--ray study of the nucleus and surrounding regions has been carried out by Stevens, Forbes \\& Norris (1998). The radio nucleus does not appear to have an X--ray counterpart. Furthermore the X--ray emission is consistent with star formation processes. Stevens \\etal conclude that if NGC 1365 harbours a black hole it is largely inactive.  \\noindent {\\bf NGC 6221} Located in a small group, NGC 6221, is a barred galaxy with a weak Seyfert nucleus. The galaxy may be interacting with NGC 6215 and has a nuclear bar (Koribalski 1996). The radio emission from NGC 6221 is extended in an symmetric bar--like structure. The spectral index of the nucleus and bar are non--thermal with $\\alpha$ $\\sim$ --0.6, indicative of SNRs. The radio morphology and other properties are all consistent with star formation being the dominant source of radio emission.  \\noindent {\\bf NGC 6810} This early type spiral is the most distant in our sample and has not been well studied to date. It may contain a bar and ring structure (Buta 1995), and does not appear to have been imaged before at radio wavelengths. Our radio image reveals a dominant nucleus surrounded by diffuse extended radio emission. Both the nucleus and surrounding region have flat spectral indices, but the brightness temperature is too low for this to be attributable to synchrotron self--absorption, which suggests that the radio spectrum is flattened by free--free absorption from young star formation. The radio morphology and other radio properties are all consistent with star formation being the dominant source of radio emission. Interestingly, recent high resolution optical spectra do not confirm the status of NGC 6810 as a Seyfert galaxy (Heisler 1998), thus it appears to have been misclassified.   \\noindent {\\bf NGC 7582} This narrow line X--ray galaxy is located in the Grus loose group along with NGC 7590 (a Seyfert 2; Ward \\etal 1980), NGC 7552 (a starburst; Forbes \\etal 1994) and NGC 7599. Several HI bridges connect group members (Koribalski 1996). Morris \\etal (1985) provide evidence for both a rapidly--rotating $\\sim$ 1 kpc ring of circumnuclear star formation and high excitation gas moving outwards from the nucleus.  Ulvestad \\& Wilson (1984) imaged the galaxy at 6 cm using the VLA with a beam size of $\\sim$ 1.5 arcsec. They measured a total 6cm flux of 69 mJy. We find a linear, double--peaked morphology to the radio emission. The southern peak appears to lie at the centre of the outer radio isophotes and is presumably the true nucleus. The nucleus has a steep spectral index ($\\alpha$ = --0.7) indicating non--thermal emission. To the NW by $\\sim$ 3 arcsec lies a second peak, which could be a second nucleus. However, it lies roughly along the bar/major axis position angle (P.A. $\\sim$ 150$^{\\circ}$). It has a 2 arcsec diameter 6cm flux of 10 mJy and a spectral index of --0.9. This second peak could therefore be a radio jet or simply a discrete star formation region occurring along the galaxy bar. The inferred SN rate in the central 6 arcsec (870 pc) is the highest in our sample (except possibly for NGC 4945) at about 1 SN every 8 years. Although the  radio morphology suggests a linear Seyfert jet, the spectral index, radio--FIR ratio, and [FeII]--radio ratio are all consistent with star formation being the dominant source of radio emission.   \\noindent {\\bf Circinus} The Circinus galaxy is perhaps the closest Seyfert galaxy known but is difficult to observe due to its proximity to the Galactic plane and large internal obscuration. Confirmation of an AGN comes from the the presence of high excitation coronal lines (Oliva \\etal 1994), X--ray emission (Matt \\etal 1996), and a compact radio core (Heisler \\etal 1998).  Like NGC 4945, Circinus is one of the few water megamaser galaxies. The megamasers in Circinus are stronger but less extreme than those in NGC 4258, and have the curious property of fluctuating on a time scale of minutes (Greenhill \\etal 1997), indicating a particularly compact source. Preliminary VLBI imaging of the megamasers (Ellingsen \\etal 1998) indicates that the maser region is extended with a velocity gradient aligned with that of the parent galaxy, and perpendicular to the jet. We regard this as strong evidence for a massive black hole in this galaxy.   Marconi \\etal (1994) found both a circumnuclear starburst ring and an ionisation cone. The [OIII] ionisation cone is asymmetric extending only to the NW, with some high excitation lines also seen in the cone region. They estimated the extinction to the nucleus to be A$_V$ $\\sim$ 20. The HI gas distribution shows spiral arms, a bar and a central `hole' (Koribalski 1996, Jones \\etal 1998). High resolution observations of the central region indicate a rapidly--rotating gas ring with a diameter of $\\sim$ 400 pc (Koribalski 1996). Observations with the ATCA have been carried out at 13 and 20 cm by Elmouttie \\etal (1995). They found extended radio lobes perpendicular to the galaxy major axis (position angle = 30$^{\\circ}$) with a spectral index of $\\alpha \\sim$ --0.7.  Our 3 and 6cm radio images, observed as part of this project, have been published, along with near--infrared line images, by Davies \\etal (1997). The radio data indicate that the nucleus is marginally resolved with a flat spectral index. The low brightness temperature indicates that this is due to free--free absorption rather than synchrotron self--absorption in a compact AGN source. There are also faint hints of extended emission which may be associated with outflowing material. Despite the clear indication of a compact AGN in the radio images, the other radio indicators are consistent with the more extended radio emission being dominated by star formation activity.\\\\"
    },
    "9804/astro-ph9804251_arXiv.txt": {
        "abstract": "The innermost regions of quasars can be resolved by a gravitational-lens {\\lq}telescope{\\rq} on scales down to a few AU. For the purpose, X-ray observations are most preferable, because X-rays originating from the innermost regions, can be selectively amplified by microlensing due to the so-called `caustic crossing'. If detected, X-ray variations will constrain the size of the X-ray emitting region down to a few AU. The maximum attainable resolution depends mainly on the monitoring intervals of lens events, which should be much shorter than the crossing time. On the basis of this idea, we performe numerical simulations of microlensing of an optically-thick, standard-type disk as well as an optically-thin, advection-dominated accretion flow (ADAF). Calculated spectral variations and light curves show distinct behaviors, depending on the photon energy. X-ray radiation which is produced in optically thin region, exhibits intensity variation over a few tens of days. In contrast, optical-UV fluxes, which are likely to come from optically thick region, exhibit more gradual light changes, which is consistent with the microlensing events so far observed in Q2237+0305.  Currently, Q2237+0305 is being monitored in the optical range at Apache Point Observatory. Simultaneous multi-wavelength observations by X-ray sattelites (e.g., ASCA, AXAF, XMM) as well as HST at the moment of a microlens event enable us to reveal an AU scale structure of the central accretion disk around black hole. ",
        "introduction": "The high power output from quasars is usually attributed to the combination of a supermassive black hole with a surrounding accretion disk. This belief is supported by a number of observations that indicate the presence of a deep gravitational potential well or a hot gas disk at the center of quasars or other active galactic nuclei; e.g., measurements of stellar velocity dispersion clearly showed a peculiar increase toward the center (Young et al. 1978; Sargent et al. 1978; see also Ford et al. 1994; Harms et al. 1994). Malkan (1983) found that the optical to UV spectra are well fitted by the standard-type accretion disk model (Shakura \\& Sunyaev 1973). Recently, by far the best evidence of a supermassive black hole has been found by radio observations of nuclear H$_2$O maser sources in NGC4258 (Miyoshi et al. 1995). Alternatively, we can infer the presence of a relativistic object from the asymmetric Fe line profile (Tanaka et al. 1995). These observational results are all attractive, but still the real vicinity of a putative black hole has not been resolved.  Q2237+0305 (e.g., Huchra et al. 1985) is the first object, in which the quasar microlensing events were detected (Corrigan et al. 1987; Houde \\& Racine 1994; see also Ostensen et al. 1996). These observations suggest that microlensing events take place roughly once per year. This rather high frequency is consistent with the microlens optical depth of $\\tau \\sim 0.8$ obtained by the realistic simulation of the lensing galaxy (i.e., Wambsganss \\& Paczy\\'nski 1994). We consider, here, specifically the microlensing events of this source caused by the so-called `caustic crossings' (see Yonehara et al. 1997 for single-lens calculations). Several authors have already analyzed this `caustic' case based on a simple model for quasar accretion disk (e.g., Wambsganss \\& Paczy\\'nski 1991; Jaroszy\\'nski, Wambsganss \\& Paczy\\'nski 1992). So far, however, only the standard-type disk, which is too cool to emit X-rays, has been considered, and thus the property of an X-ray microlensing of quasar, e.g., Q2237+0305, has not been predicted. We stress here the significance of X-ray observations to elucidate the physics of the innermost parts of the disk, since X-rays specifically originate from a deep potential well. The observations allow us to assess the extension of hot regions on several AU scales and resultantly to deduce the mass of a central massive black hole.  In this $Letter$, we propose to investigate quasar central structure by using X-ray microlensing of Q2237+0305. In section 2, we describe the method for resolving X-ray emission properties of the inner disk structure on a scale down to a few AU. In section 3, we calculate the disk spectra and light curves during microlensing. We here use realistic disk models: the optically-thick, standard disk (Shakura \\& Sunyaev 1973) and the optically-thin, advection-dominated accretion flow (ADAF, Abramowicz et al. 1995; Narayan \\& Yi 1995; see also Ichimaru 1977). ",
        "conclusions": ""
    },
    "9804/cond-mat9804137_arXiv.txt": {
        "abstract": "We study with Monte Carlo methods an ensemble of $c=-5$ gravity graphs, generated by coupling a conformal field theory with central charge $c=-5$ to two-dimensional quantum gravity.  We measure the fractal properties of the ensemble, such  as the string susceptibility exponent $\\gamma_s$ and the intrinsic fractal dimensions $d_H$.  We find $\\gamma_s = -1.5(1)$ and $d_H = 3.36(4)$, in reasonable agreement  with theoretical predictions.  In addition, we study the critical behavior of an Ising model on a {\\it quenched} ensemble of the \\mbox{$c=-5$} graphs and show that it agrees, within numerical accuracy, with theoretical predictions for the critical behavior of an Ising model coupled {\\it dynamically} to two-dimensional quantum gravity, provided the total central charge of the matter sector is $c=-5$.  From this we conjecture that the critical behavior of the Ising model is determined solely by the average fractal properties of the graphs, the coupling to the geometry not playing an important role. ",
        "introduction": "Randomness in statistical systems arises in a variety of situations and is a very rich and complex subject. Quenched randomness is frequently used in studying the role of impurities and inhomogeneities in real physical systems where the characteristic time-scale of the disorder is much longer than other dynamics of the system.  Annealed randomness, on the other hand, arises naturally in studies of fluctuating geometries, such as two-dimensional quantum gravity or fluid membranes, where the disorder is dynamically modified by interaction between the geometry and matter fields living on the surfaces.  For a statistical system coupled to random disorder, either in a quenched or annealed approach, the main question is to assess the effect randomness has on the critical behavior of the pure system.  One prediction in this direction is the Harris conjecture \\cite{harris} which states that randomness changes the values of critical exponents only if the specific heat exponent $\\alpha$ of the pure system is positive. This conjecture has been studied in many models with quenched disorder, such as the $2d$ Ising model \\cite{dots} (where the Harris criterion is ambiguous as $\\alpha = 0$) and the Potts model \\cite{bondpott}. For both models a change in the critical behavior is observed.  All the above mentioned studies deal with weak disorder. More recently the critical behavior of systems on lattices with fractal structure very different from a flat surface has been investigated.  Such systems arise naturally when matter, in the form of conformal field theories, is coupled to two-dimensional quantum gravity.  These models can be studied either in a continuum formulation, by Liouville field theory, or using discretized approaches like, for example, models of dynamical triangulations, formulated either as matrix models or studied with numerical simulations. For these systems the disorder is, however, different from the one discussed above in that it is annealed, i.e.\\ the models couple dynamically to fluctuations in the geometry.  A remarkable degree of universality does emerge for models coupled to two-dimensional quantum gravity. Namely, the change in the critical behavior of the systems, and their back-reaction on the geometry, only depends on the total central charge of the matter sector.  This manifests itself in the so-called KPZ scaling relation which describe how the scaling dimensions of conformal operators are changed by the interaction with gravity \\cite{kpz}.  Moreover, this universality also extents to the fractal structure of the surfaces, from which we derive the string susceptibility exponent $\\gamma_s$ and the fractal dimension $d_H$.  In view of this universality it is tempting to conjecture that the critical behavior of a particular system, when coupled to a fluctuating geometry, only depends on the (average) fractal structure of the surface. Details of the interaction between the system and the geometry, or the geometrical fluctuations, are not important as such --- they only serve the purpose of defining the average fractal geometry.  If this conjecture is true it implies that how the average over disorder is performed, i.e.\\ that the disorder is annealed, is not essential.  In particular, predictions of the KPZ scaling relation for the change in the critical behavior should just as well apply to  models with quenched disorder, {\\it provided the quenched average is taken over the same ensemble of disorder as is generated in the annealed approach}.  There are some recent simulations that have addressed the question of the critical behavior of spin models on a quenched ensemble of graphs generated by two-dimensional quantum gravity. Both the Ising model \\cite{bhj} and the 10-state Potts model \\cite{bjj} have been studied on an ensemble of pure gravity graphs ($c=0$).  For the Ising model a critical behavior compatible with an Ising model coupled dynamically to gravity was found, although the accuracy of the results is not sufficient to rule out the conjecture discussed above.  The goal of this paper is two-fold.  First, we want to investigate the fractal geometry of two-dimensional quantum gravity coupled to a conformal field theory with central charge $c=-5$.  More precisely, we want to determine the fractal dimension of the corresponding surfaces, using recently developed finite-size scaling methods \\cite{hausd,janhaus} and to compare it to the (contradictory) theoretical predictions that exist \\cite{anhaus1,anhaus2}.  Second, we want to investigate the critical behavior of an Ising model on a quenched ensemble of $c=-5$ graphs and to compare it with predictions from Liouville theory, for the critical behavior of an Ising model coupled dynamically to two-dimensional quantum gravity, and to verify, or disprove, our conjecture about the effect of the disorder.  Our motivation for choosing $c=-5$ is that both its predicted fractal structure and the critical behavior of the Ising model is substantially different from both a flat space and for a pure two-dimensional quantum gravity.  This makes these different critical behavior easier to distinguish in numerical simulations.  The paper is organized as follows: In Section~2 we study the fractal properties of a $c=-5$ conformal field theory coupled to two-dimensional gravity.  We define the model in Section~2.1 and discuss the details of the simulations in Section~2.2. In Sections~2.3 and 2.4 we present our measurements of the string susceptibility exponent $\\gamma_s$ and of the fractal dimension $d_H$.  And in Section~2.5 we comment on how this particular ensemble of graphs differs from other types of graphs frequently used in studying disordered system. The second part of the paper deals with an Ising model on the $c=-5$ graphs in a quenched approach. In Section~3.1 we discuss the prediction from Liouville theory for the critical behavior of an Ising model coupled dynamically to two-dimensional gravity. In Section~3.2 we discuss details of the simulations and the observables we use to probe the critical behavior. In Sections~3.3 and 3.4 we determine the critical temperature of the Ising model and the corresponding critical exponents. Finally, in Section~4 we summarize and discuss our results. ",
        "conclusions": "The main results of the work presented in this paper can be summarized as follows: \\begin{itemize} \\item[({\\it a})] The fractal dimension of  surfaces, defined by a conformal field theory with central charge $c=-5$ coupled to two-dimensional quantum gravity, is $d_H = 3.36(4)$.  This is in reasonable agreement with, and supports, the theoretical prediction Eq.~(\\ref{andH2}), whereas it definitely rules out Eq.~(\\ref{andH1}). \\item[({\\it b})] The critical behavior of an Ising model on a {\\it quenched} ensemble of $c=-5$ graphs agrees well with the predictions, from the KPZ scaling relation, for an Ising model on an {\\it annealed} ensemble of graphs with {\\it identical} fractal properties. \\end{itemize}  The first result, especially combined with the recent simulations of $2d$--gravity for $c=-2$ \\cite{cm2}, lends a strong support to Eq.~(\\ref{andH2}) as a correct description of the fractal structure of two-dimensional quantum gravity for $c \\leq 0$.  This makes, however, its disagreement with numerical simulations in the region $0 < c \\leq 1$ all the more surprising.  What is it in derivation of Eq.~(\\ref{andH2}) that breaks down for $c>0$? Or are the simulations dominated by finite-size errors and simulations of larger systems will eventually agree with Eq.~(\\ref{andH2})?  The result for the Ising model is even more interesting. As the theoretical predictions are obtained for an Ising model coupled dynamically to the disorder, this supports the conjecture put forward in the Introduction about the equivalence between annealed and quenched averages over disorder. That is, the only thing relevant for the critical behavior of the Ising model are the average fractal properties of graphs the spins ``see''. How the statistical average over graphs is performed, quenched or annealed, is not relevant.  It is also worth noting that we can continuously change the average fractal properties of the graphs by changing the embedding dimension $D$ in Eq.~(\\ref{eq213}). This allows a continuous interpolation between a flat surface and surfaces corresponding to pure gravity. If the prediction of Liouville theory, the KPZ formula, holds for all those models, this implies that the critical behavior of the Ising model should change continuously in the process. In the language of the renormalization group this implies a continuous line of fixed points, rather than isolated points.  There are well known examples of this; the low-temperature phase of the two-dimensional $XY$--model or the critical line of the Ashkin-Teller model. But is this statement also true for very weak disorder? If we change the fractal dimension infinitesimally, from 2 to $2+\\epsilon$, is that  enough to change the critical behavior of the Ising model?  Or, alternatively, does there exist some central charge $c^{\\prime} < -5$ were the geometrical disorder is not strong enough and we always get the Onsager exponents? This point deserves further study.  One could also look at the examples of weak disorder that have been studied recently, for example the site or bond-diluted Ising model, and ask if that kind of disorder can also be classified according to some average fractal properties of the lattices. And, moreover, if one could observe some kind of universality in the critical behavior, depending on the fractal structure, akin to what we have presented in this paper.  In view of how dramatically the critical behavior of the Ising model changes on surfaces with such strong disorder, one might ask if such change could be observed in real physical systems. Possible candidates for such systems could be, for example, electrons trapped on the interfaces between two liquids, or on the surface of some porous material, were the surfaces had some well defined non-trivial fractal structure. As our results indicate, it is only the average geometry of the surfaces that is important for the Ising model, not its fluid nature or curvature fluctuations.  Thus the relative time-scale between the interactions of the particles and the change in the geometry should be irrelevant.  \\vspace{20pt} \\noindent {\\bf Acknowledgments:} The work of G.T.\\ was supported by the Humboldt Foundation. The work of P.B.\\ was partially supported by KBN grants 2P03 B19609 and 2P03 B04412."
    },
    "9804/astro-ph9804100_arXiv.txt": {
        "abstract": "I describe a general framework that could allow to understand the broad band spectra  of blazars and lead to a unified picture of the emission from relativistic jets, in BL Lac objects as well as in flat spectrum, radio loud Quasars. The scheme serves as a useful basis to introduce and discuss some of the most interesting results so far obtained on Blazars with {\\it Beppo}SAX. ",
        "introduction": "The \"blazar phenomenon\" is due to the presence of  relativistic flows (jets) emanating from the nuclei of active galaxies which are radio loud. The power to energize the radio lobes is transported in the jets. Blazars are the subset of radio loud AGN  for which the  relativistic jet happens to point at small angles to the line of sight. Since the radiation emitted by the jet is relativistically beamed into a narrow cone along the direction of motion,  the aligned observer will receive a strongly enhanced flux. For a bulk Lorentz factor $\\Gamma\\simeq 10$ at an angle $\\theta\\simeq 1/\\Gamma$ the flux enhancement factor is $10^3-10^4$.  The evidence in favor of this picture has been accumulating and is now solid (e.g. \\cite{pu} and refs therein) although the origin of the jets is poorly understood and their physical parameters are highly uncertain. The relativistically amplified,  non thermal emission from high energy particles in the jet can account for the extreme properties of blazars concerning variability, polarization and energy distribution of the continuum, which extends from the radio to the gamma-ray band.  Traditionally BL Lac objects, where no prominent emission lines are observed (with an upper limit of 5 \\AA), were thought to represent a separate class perhaps more extreme than Quasar-like blazars. The latter include optically violently variable and highly polarized quasars (OVV, HPQ) or more generally Quasars with flat radio spectrum (FSRQ) indicating strong emission from the self absorbed core. It has become clear however that BL Lacs have on average lower luminosity than quasar like blazars (\\cite{pado}) and that the distribution of emission line equivalent widths is  continuous (\\cite{scafal}). We will therefore in the following consider blazars as a single class of objects, implicitly assuming that, irrespective of the emission line properties which derive from the surrounding gas, the same physical mechanisms operate in  relativistic jets over a wide range of luminosities.  By studying the blazar continuum we expect to learn about the radiation mechanisms in the jets, about the processes of particle acceleration and energy transport along the jets and ultimately about their origin and evolution.  \\begin{figure*}[bt] \\vspace{9pt} \\psfig{file=sed_medie.ps,width=15.0truecm,height=11.5truecm,rheight=8.7truecm} \\caption{\\small\\sf Average SEDs for the ``total blazar sample'' binned according to radio luminosity irrespective of the original classification.  The overlayed dashed curves are analytic approximations obtained assuming that the ratio of the peak frequencies is constant and that the luminosity at the second peak is proportional to the radio luminosity (from \\protect\\cite{foss98} } \\label{fig:sed_medie} \\end{figure*}  \\section {The broad band spectra of blazars}  The discovery by EGRET on board CGRO of copious $\\gamma$-ray emission from blazars caused a \"Renaissance\" in this field. In fact it had been noted early on that the high density of relativistic electrons necessary to produce the observed compact synchrotron emission would lead to strong, even catastrophic inverse Compton radiation (\\cite{hbs}).  X-ray measurements were used to constrain the amount of inverse Compton emission allowed and to derive minimum values for the necessary beaming factors (e.g. \\cite{gp93}). At present the $\\gamma$-ray observations allow to measure the intensity and spectral shape of a component which contains a substantial fraction and in some cases the bulk of the emitted power leading to strong constraints on the physical parameters of the emitting region.  It is clearly important, besides observing single objects, to try to derive general properties of the continuum and understand whether  and how they differ for instance in BL Lac objects and FSRQ. It is especially interesting to discuss whether the gamma-ray emission is a general property of the whole class. We have recently addressed this problem (\\cite{foss98}) collecting multifrequency data for three complete samples of blazars: the 2 Jy sample of FSRQs, the 1Jy sample of BL Lac objects and the sample of BL Lacs selected in the X-ray band from the Einstein Slew Survey. Systematic differences in the shape of the continuum in specific spectral bands among different subclasses of blazars were noted early on (e.g. \\cite{gg86,imp,ww,smu,umu}. In particular we note that the percentage of objects detected with EGRET (100 MeV - 10 GeV) is significantly larger for the sample of FSRQ than for the two BL Lac samples (40 \\% vs. 26\\% and 17\\% respectively. Nevertheless plotting the average SEDs as shown in Fig.~\\ref{fig:sed_medie}, we can see that the shapes are \"globally similar\".  In Fig.~\\ref{fig:sed_medie} all blazars in the three complete samples were merged and  grouped in luminosity classes irrespective of their original classification and the dashed lines drawn for comparison derive from an analytic parametric representation (\\cite{foss98}).  The main results of this work are the following: \\begin{itemize} \\item two peaks are present in all the SEDs  \\item the first peak occurs at lower frequencies for the highest luminosity objects  \\item the frequency at which the second peak occurs correlates with that of the first one. The dashed curves correspond to a constant ratio between the two peak frequencies.  \\end{itemize}  For the most luminous objects the first peak is at frequencies lower than the optical band while for the least luminous ones the reverse is true. Thus highly luminous objects have a \"red\" (steep) IR to UV continuum while objects of lower luminosity have a bluer IR to UV continuum. For this reason and to recall intuitively the location of the peaks on the frequency axis we will briefly call \"red\" blazars the objects in the three highest luminosity classes and \"blue\" blazars those in the two lower luminosity classes. The present data suggest a continuous spectral sequence and no absolute separation between red and blue blazars.  Considering the continuum of different objects in a fixed spectral range its shape changes systematically with luminosity along the sequence, as the peak frequency approaches and  moves across the chosen frequency interval. In particular the X-ray spectrum becomes steeper and the gamma-ray spectrum (in the EGRET range) becomes flatter from \"red\" to \"blue\" blazars, as the two peaks march to higher frequencies. The different location of the gamma-ray peak can account for the different detection rates of BL Lacs and FSRQs by EGRET. Objects whose $\\gamma$-ray emission peaks in the EGRET range are more easily detected.  Recently ground based observations in the TeV range performed with Cherenkov telescopes have detected two of the X-ray brightest \"blue\" blazars (refs). We expect that with the progress in sensitivity  many more will be detected giving access to the study of the highest energies from ground.    A final comment concerns variability. It is interesting to note that the largest variability is usually observed close to or above each of the two peaks and is usually in the sense of a hardening of the spectrum at higher intensity. These statements are  based mostly on observations at high energies (X-rays and gamma-rays) and concern a limited number of sources (e.g. \\cite{umu}) therefore they should be considered as tentative suggestions rather than established facts.  \\section {Interpretation}  It is generally thought that the first spectral component peaking at far infrared up to X-ray frequencies is due to synchrotron radiation. The spectra from the radio to the submm range most likely involve contributions from different regions of the jet with different self absorption turnovers. However, from infrared frequencies upwards the synchrotron emission should be thin and could be produced essentially in a single homogeneous region. Inverse Compton scattering of soft photons by the high energy electrons emitting the thin synchrotron radiation could be responsible for the second ( high frequency) component of the SED, peaking in the gamma-ray band. The soft photons could be the synchrotron photons themselves (SSC) or photons outside the jet (EC), possibly produced by an accretion disk or torus and scattered or reprocessed by the surrounding gas (e.g \\cite{sbr}, \\cite{umu} and refs therein).  If the same region is responsible for the two spectral components then, irrespective of the nature of the seed photons, {\\it the two peaks must derive from  the same high energy electrons}. Therefore a change in the density and/or spectrum of those electrons is expected to cause correlated variability at frequencies close to the two peaks. In the SSC model the inverse Compton intensity is expected to vary more than the synchrotron one, approximately as the square of it in the simplest case while in the EC model one expects a linear relation (\\cite{miami}).  Measuring the two peaks simultaneously is thus the best means to determine the physical parameters of the emission region and studying the variability of the spectra around the peaks yields unique insight into the mechanisms of particle acceleration and  energy loss in the jet. The variability correlation should enable to disentangle the contribution of different sources of seed photons (SSC vs. EC).  The \"spectral sequence\" discussed above could be attributed to a systematic dependence of the critical electron energy (the break energy) and/or of the magnetic field on luminosity. Assuming that the beaming factors are not significantly different along the sequence, the trend in apparent luminosity is also a trend in intrinsic luminosity. In the SSC model the break energy of the electrons is univocally determined by the ratio of the frequencies of the two peaks and should therefore be approximately constant. \"Red\" blazars should then have lower magnetic field than \"blue\" blazars. Systematic model fitting of all the $\\gamma$-ray detected blazars with sufficient multifrequency data suggest  that as the magnetic  energy density decreases the external photon energy density becomes important so that a smooth transition between the SSC and EC scenario takes place (\\cite{gg98}).       \\section {SAX observations}  The X-ray band is crucial for a discussion of the above problems in that the synchrotron and inverse Compton components which have different spectral shapes may both be relevant. Simultaneous observations over a broad energy range are required to disentangle the two mechanisms. The  unique characteristics of the {\\it Beppo}SAX instrumentation  appear therefore ideal for blazar studies. Observations of bright blazars detected in $\\gamma$--rays were proposed with the main aims of:  \\begin{itemize} \\item determining the spectral shape up to the 100 keV range, thus exploring the connection between X-rays and gamma-rays  \\item studying the variability in relation with other wavebands, especially $\\gamma$--rays. \\end{itemize}  \\begin{figure*}[bt] \\vspace{9pt} \\psfig{file=sed_3c279.ps,width=15.0truecm,height=13truecm,rheight=11.2truecm} \\caption{\\small\\sf Spectral energy distribution during the 1997 campaign, compared with previous ones. The 1997 data are from: X--rays = {\\it Beppo}SAX data; $\\gamma$--rays = Hartman, private comm.; R-band = Raiteri \\& Villata, private comm.} \\label{fig:sed_3c279} \\end{figure*}  In the following I will briefly mention and comment some of the most interesting results obtained so far. Besides 3C 273 (\\cite{htv}) which is probably intermediate between a blazar and a \"normal\" quasar,  two flat spectrum radio quasars, 3C 279 and PKS 0528+134, were observed (before June 1997) and found in a low intensity state.  In the scheme presented above these are \"red\" blazars. I will discuss here the first source (also \\cite{mtv}), while for the second one I refer to \\cite{geltv}. Finally I will consider results on \"blue\" blazars ( Mrk 421, 1ES 2344+514 and Mrk 501 (see the contributions in this volume \\cite{ftv}, \\cite{gitv} and \\cite{gtv} respectively).  It is important to remember that in \"red\" blazars the X-ray emission represents the lower energy end of the inverse Compton emission, while for \"blue\" blazars it is the high energy end of the synchrotron emission.   \\subsection {3C 279}   The X-ray spectrum measured with {\\it Beppo}SAX in January 1997, the simultaneously measured optical flux  and the quasi simultaneous gamma-ray flux (Hartman, private communication) are shown in Fig.~\\ref{fig:sed_3c279} together with other simultaneously measured SEDs obtained at other epochs: the high state observed in June 1991 (\\cite{hart}), the low state observed in January 1993 (\\cite{m94}) and the preflare and flare states observed in January - February 1996 (\\cite{wehrle98}).  At the epoch of the {\\it Beppo}SAX observations the  $\\gamma$--ray flux was a factor 6 and 20 weaker respectively than  measured in June 1991 and early February 1996. In X-rays the amplitude is smaller but there is a good correlation between the X-ray and gamma-ray fluxes especially in the 2-10 keV band (see Fig. 3 in Maraschi et al. this volume). Note that the fluxes at 1 keV measured by {\\it ROSAT} and {\\it Beppo}SAX  in the 1993 and 1997 low states are similar   However the spectrum measured by {\\it Beppo}SAX  is significantly flatter providing a good connection with the higher $\\gamma$--ray flux in 1997. The simultaneity (within one day) of the X-ray (XTE) and  $\\gamma$--ray peaks during the 1996 flare suggests that the X-ray to gamma-ray emission originates in a single region and that the spectrum hardens with increasing intensity. It is possible  that the $\\gamma$-ray peak in the SED  moves to higher energies at the flare peak.  The situation  is much more complex at lower energies. Although there is still a general correlation of the IR-optical-UV fluxes with the gamma-ray intensity on long timescales, the flux variation at optical wavelengths corresponding to the rapid 1996 flare is quite small. Note also that for 3C 279 the (presumed) peak of the synchrotron component falls in an unexplored region of the spectrum, between $10^{12}$ and $10^{14}$ Hz.  In the SSC model one expects that the inverse Compton emission varies with the square of the amplitude of the synchrotron emission due to the same electrons (e.g. at the two peaks of the SED). This is compatible with the long term variations but not with the strong rapid flare observed in 1996, where the amplitude in gamma-rays was larger than the square of the optical one. On the other hand if the seed photons for the inverse Compton process are external to the jet, they should not be rapidly variable and the inverse Compton emission is expected to vary linearly with the synchrotron one. Thus neither of the two \"simple\" models can adequately account for the multifrequency variability behaviour.  A possible way out is that the seed photons derive from backscattering and/or reprocessing of radiation produced in the jet by gas clouds closely approaching the jet itself (\\cite{ggmadau}). This model is attractive and needs to be studied in more detail. Another possible way out is that the region emitting the synchrotron radiation is inhomogeneous so that the observed variability is diluted by a more stationary component.  \\subsection {Blue blazars}  The X-ray emission from these objects has been observed to vary dramatically on short timescales at least in the brightest prototypes, PKS 2155-304 and Mrk 421. This can be understood recalling that the X-ray emission represents the high energy end of the synchrotron component: it is therefore due to radiation from the highest energy electrons which have the shortest lifetimes and can vary very rapidly.  A continuous acceleration mechanism must be responsible for maintaining in the source particles which have lifetimes of the order of hours. However the injection mechanism may be continuous only in an average sense or at a low intensity level and  episodes of increased injection rate may occur causing variability. The study of flares and of spectral variability associated with them  gives direct information on the spectra of the freshly injected/accelerated particles and their subsequent decay to a state of quasi-equilibrium.  The photons upscattered through the inverse Compton process by the highest energy electrons reach  TeV energies so that the \"bluest\" and brightest sources can be detected from ground based Cherenkov telescopes. This is the case up to now for three objects: Mrk 421, 1ES 2344+514 and Mrk 501. All of them were observed with SAX in the first half of AO1. The simultaneous observation of a source in X-rays and at TeV energies should allow to determine unambiguously the energy of the radiating electrons, the beaming factor, the magnetic field and the energy density of the seed photons. In some cases it may be necessary to take into account that scattering will occur in the Klein Nishina regime.  \\begin{figure*}[bt] \\vspace{9pt} \\psfig{file=sed_mkn421.ps,width=15.0truecm,angle=270,height=13truecm,rheight=11.2truecm} \\caption{\\small\\sf The May 97 data from {\\it Beppo}SAX are compared with the  energy distributions measured in 1994 (\\protect\\cite{macomb}} \\label{fig:sed_mkn421} \\end{figure*}  \\subsection {Mrk 421}  This source has been repeatedly observed with ASCA and at other relevant wavelengths including TeV observations. (\\cite{macomb})  The {\\it Beppo}SAX observations of Mrk 421 (Fossati et al. this volume) show a decay between two intensity states closely similar to those  previously observed with ASCA. The higher state shows a flatter spectrum indicating that the emission peaks at higher energies. In Fig.~\\ref{fig:sed_mkn421} the {\\it Beppo}SAX data are compared with the ASCA and TeV data presented by \\cite{macomb}.  Mastichiadis and Kirk \\cite{mk}have shown that this behaviour could be the result of an injection mechanism in which the maximum energy of the injected particles increases. The two \"states\" would therefore represent equilibrium states for injection spectra extending to different maximum energies. The fact that the spectral variability observed in 1994 is closely reproduced in the  {\\it Beppo}SAX observations indicates that the involved region is the same (same physical parameters). In addition SAX detected the source at higher energy with the PDS. The preliminary analysis yields a flat spectrum at high energies suggesting that the inverse Compton component becomes dominant in this band. The variation in the PDS is much smaller than in the MECS which is consistent with attributing the hard emission to IC from electrons of much lower energies. An extensive campaign for simultaneous ASCA and TeV observations will take place in  the spring of 1998, to which {\\it Beppo}SAX observations could add significantly.   \\subsection {1ES 2344+514, Mrk 501}  This source is not as well studied as Mrk 421 but  shows a  similar, more extreme behaviour (\\cite{gitv}). The emission peak was in the medium X-ray range in December 96 and shifted to 20-50 keV during the flare of a factor 2 observed between December 3-7 96.  The most extraordinary spectral variation was found in Mrk 501 (\\cite{pian98}, \\cite{gtv}). The source was in a state of strong activity at TeV energies although the 1 keV flux had not increased dramatically. However compared to previous observations {\\it the X-ray spectrum had changed dramatically} being flatter than or close to 1 up to 100 keV. In correspondence to a large TeV flare , the 1 to 100 keV spectrum hardened still indicating that the emission peak was at or above 100 keV, while past multifrequency measurements  showed it to be below 1 keV. Thus in Mrk 501 the shift of the peak frequency was more than a factor 100. Some theoretical implications are discussed by Ghisellini (\\cite{gtv}).  This unprecedented behaviour may be less uncommon than judged at first sight. In fact in the medium X-ray band the intensity behaviour was not exceedingly dramatic and good spectral capabilities up to the hard X-ray band, necessary to reveal the phenomenon, have only recently become available with {\\it Beppo}SAX.  It is interesting to mention that a spectral survey of BL Lacs selected in relatively soft X-rays yields evidence of some hard X-ray spectra (\\cite{wolter}). Sources which have more or less permanent emission peaks in the hard X-ray range may also exist and have gone undetected so far. ",
        "conclusions": "The results  from the first part of the {\\it Beppo}SAX  AO1 Core Program on Blazars are extremely exciting. Bright $\\gamma$-ray blazars can be usually detected up to 50 - 100 keV allowing detailed study of the synchrotron and inverse Compton emission and of the correlations between their temporal and spectral variability.   The observations  presented above suggest and support:  \\begin{itemize} \\item the correlation between the X--ray (1 keV) and $\\gamma$--ray (0.1-10 GeV) fluxes, which now holds over a period of 6 years and about a factor 30 in flux change for $\\gamma$-rays (3C 279 -- PKS 0528 + 154)  \\item the smaller amplitude of variability of the synchrotron component compared to the high energy one.  \\item the synchrotron emission peak seems to shift systematically to higher energies during flares (Mrk 421, 1ES 2344+541, Mrk 501)  \\end{itemize}  These results will undoubtedly stimulate new observations and new theoretical approaches.  In particular genuinely time dependent models  for the acceleration of particles and the spectral evolution of the emitted radiation are needed."
    },
    "9804/astro-ph9804336_arXiv.txt": {
        "abstract": "In this paper we present a detailed study of the radio galaxy J1324$-$3138, located at a projected distance of 2$^{\\prime}$ from the centre of the Abell cluster of galaxies A3556, belonging to the core of the Shapley Concentration, at an average redshift z=0.05. We have observed J1324$-$3138 over a wide range of frequencies: at 327 MHz (VLA), at 843 MHz (MOST), and at 1376 MHz, 2382 MHz, 4790 MHz and 8640 MHz (ATCA).  Our analysis suggests that J1324$-$3138 is a remnant of a tailed radio galaxy, in which the nuclear engine has switched off and the radio source is now at a late stage of its evolution, confined by the intracluster gas. The radio galaxy is not in pressure equilibrium with the external medium, as it is often found for extended radio sources in clusters of galaxies. We favour the hypothesis that the lack of observed polarised radio emission in the source is due to Faraday rotation by a foreground screen, i.e. the source is seen through a dense cluster gas, characterised by a random magnetic field.  An implication of the head-tail nature of the source is that J1324$-$3138 is moving away from the core of A3556 and that possibly a major merging event between the core of A3556 and the subgroup hosting J1324$-$3138 has already taken place. ",
        "introduction": "Extended radio emission associated with galaxies in clusters is often characterised by morphologies which reflect the interaction between the radio emitting plasma and the local environment in the cluster. Head-tail sources are usually associated with non-dominant cluster galaxies moving at a considerable speed within the cluster. Their morphologies are then explained as the result of ram pressure exerted by the intergalactic medium on the double sided radio emission (see for example O'Dea \\& Owen 1985a and 1985b, and Owen 1996, for a recent review). Wide-angle tail radio galaxies, on the other hand, are more difficult to account for with the above mentioned model, since they are usually associated with massive and dominant cluster galaxies, with much lower peculiar velocities with respect to the cluster mean. Beyond ram pressure, it is now accepted that large flows of hot gas could provide a wind within clusters of galaxies able to bend straight jets into wide-angle tail morphologies (Owen 1996 and references therein).  \\noindent The study of extended galaxies in clusters is important for a variety of reasons. The morphology and the direction of the extension may give important information on the dynamics of the galaxy, such as, for example, the direction of the motion projected on the plane of the sky. Furthermore, the non-thermal pressure in the tails of radio emission can be compared to the thermal pressure exerted by the intracluster gas, in those cases where X-ray data are available to provide estimates of the temperature and pressure. This is crucial for studying the interaction between the radio emission and the external gas, and for deriving information on the evolution of radio sources in clusters of galaxies, as well as the influence of the cluster dynamics (such as, for example, merging processes) on the radio properties of the cluster. Last but not least, the observed polarisation properties of the radio emission may give information on the intracluster magnetic field and its structure.  \\medskip In this paper we present a detailed study of the extended radio galaxy J1324$-$3138 (RA$_{J2000} = 13^h24^m01^s$, DEC$_{J2000} = -31^{\\circ}38^{\\prime}$), located in the central region of the Abell cluster A3556. It was first observed  in a radio survey of the clusters of galaxies in the Shapley Concentration core carried out at 843 MHz with the Molonglo Observatory Synthesis Telescope (MOST) and at 1376 MHz with the Australia Telescope Compact Array (ATCA) (Venturi et al. 1997, hereinafter Paper I). This work is part of a larger project whose aim is to study the radio/optical properties of the clusters in the core of the Shapley Concentration, in particular the chain formed by A3556-A3558-A3562 (Venturi et al. 1998), both from a statistical point of view and through a detailed analysis of the physical properties of the extended radio galaxies in these clusters.  In Figure 1 the superposition of the radio isophotes on the Digitised Sky Survey shows that the radio component located at the south west extremity of the extended radio emission is coincident with the nucleus of the 15.6 magnitude galaxy \\#5975 in the COSMOS catalogue (RA$_{J2000}$ = $13^h23^m57.5^s$, DEC$_{J2000}$ = $-31^{\\circ}38^{\\prime}45^{\\prime\\prime}$). Its radial velocity velocity v = 15054 km s$^{-1}$ (Stein 1996) establishes that it belongs to A3556 ($<v> = 14357$ km s$^{-1}$, Bardelli et al. 1998). This, coupled with the fact that only a few very faint optical objects fall within the envelope of the radio emission, led us to the conclusion that J1324$-$3138 is an extended, possibly head-tail, radio galaxy located in the vicinity of the cluster centre (see Paper I).  In Section 2 we present the observational data. In Section 3 the morphology of the source is described and analysed, and in Section 4 a detailed study of the synchrotron spectrum is carried out. The nature of the source, its relation to the cluster of galaxies A3556 and its implications for cluster merging and formation is discussed in Section 5.  Throughout the paper we use a Hubble constant of H$_0$ = 100 km s$^{-1}$Mpc$^{-1}$. At the redshift of the cluster this implies that 1$^{\\prime\\prime}$ = 0.67 kpc. \\begin{figure} \\epsfysize=8.5cm \\epsfbox{FIG1.PS} \\caption[]{ 4790 MHZ radio isophotes of the extended radio galaxy J1324$-$3138 and of the nearby radio galaxy J1324$-$3140, associated with the dominant cD galaxy in the centre of A3556 (see Paper I), superimposed on the DSS optical image. The resolution of the image is $20^{\\prime\\prime} \\times 10^{\\prime\\prime}$, p.a. $0^{\\circ}$.} \\end{figure} \\vskip 1.0truecm \\noindent ",
        "conclusions": "We have presented observations of the radio galaxy J1324$-$3138, located in the central region of the Abell cluster A3556, over a wide range of frequencies and resolutions. We can briefly summarise our results as follows:  \\noindent {\\it (a)} J1324$-$3138 is an example of a {\\it remnant} of a radio galaxy, i.e. a source in which the engine of the radio emission has switched off. The evolution of the radio emission is presently dominated by synchrotron losses.  \\noindent {\\it (b)} The radio source is not in pressure equilibrium with the intracluster gas. In particular it is underpressured.  \\noindent {\\it (c)} We suggest that the lack of polarisation in the source is due to the presence of an intervening Faraday screen, i.e. a cluster scale magnetised medium, as it is now often observed in clusters of galaxies, which depolarises the radio emission.  \\noindent {\\it (d)} Under the hypotesis that cluster mergers influence the radio emission of a galaxy, the properties of J1324$-$3138, coupled with the peculiarities of A3556 at radio and optical wavelengths (Paper I and Bardelli et al. 1998), suggest that the core of A3556 and the subgroup hosting J1324$-$3138 have already undergone a major merging event.  \\vskip 1.0truecm \\noindent {\\bf Acknowledgments}  We wish to thank D. Dallacasa for his suggestions and discussion while this work was carried out, and R. Fanti for careful reading of the manuscript. We are grateful to S. Ettori for providing unpublished results, and to P. Stein for giving us the spectrum of J1324$-$3138. T.V. acknowledges the receipt of two grants from CNR/CSIRO (Prot. n. 119816 and Prot. n. 088864).  The MOST is operated by the University of Sydney, with support from the Australian Research Council. The Australia Telescope Compact Array is operated by the CSIRO Australia Telescope National Facility. The National Radio Astronomy Observatory (NRAO) is operated by Associated Universities, Inc., under contract with the National Science Foundation. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration."
    },
    "9804/astro-ph9804046_arXiv.txt": {
        "abstract": "The CAT (Cherenkov Array at Th\\'emis) imaging telescope, equipped with a very-high-definition camera (546 fast phototubes with $0.12^{\\circ}$ spacing surrounded by 54 larger tubes in two guard rings) started operation in Autumn 1996 on the site of the former solar plant Th\\'emis (France). Using the atmospheric Cherenkov technique, it detects and identifies very high energy $\\gamma$-rays in the range $250\\:{\\mathrm GeV}$ to a few tens of TeV. The instrument, which has detected three sources (Crab nebula, Markarian 421 and Markarian 501), is described in detail. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804270_arXiv.txt": {
        "abstract": "We report results of a \\rosat\\ High-Resolution Imager (HRI) observation of the X-ray error box given by the \\sax\\ Wide Field Camera for the gamma-ray burst that occurred on 1997 February 28. The observation started 10 days after the burst and ended three days later, with a total exposure of 34.3~ks. An X-ray source was detected within the 3$'$ WFC error box and its position determined with a 10$''$ radius accuracy. The source position is in the \\sax\\ Narrow Field Instrument source error box and is coincident (to within 2$''$) with the optical transient associated with GRB970228. This is the most precise position obtained for an X-ray afterglow and confirms that the X-ray and optical afterglows have the same origin. We present the 0.1--2.4~keV combined HRI and \\sax\\ Low-Energy Concentrator Spectrometer decay light curve which can be well fit with a power-law. The decay is consistent with that measured at higher energies (2--10~keV) with the \\sax\\ Medium-Energy Concentrator Spectrometer. ",
        "introduction": "Observations of celestial Gamma-Ray Bursts (GRB) performed over the last 25 years had not, until recently, succeeded in finding counterparts in other wavelength bands. The ability of the \\sax\\ satellite to provide arc minute precision positions (Piro et al. 1998) and to observe these positions within hours of the GRB changed this situation in 1997 when the X-ray afterglow of GRB970228 was measured (Costa et al 1997a). The burst was detected (Costa et al. 1997a) with the Gamma-Ray Burst Monitor (GRBM) (40--70~keV, Frontera et al. 1997a) on 1997 February 28.123620 UT and also detected in the 1.5--26 keV energy range by one of the two Wide Field Cameras (WFC No. 1) aboard the same satellite (Jager et al. 1997). Its position was determined with an error circle of 3~arcmin (3$\\sigma$) radius, centered on $\\alpha_{2000}\\,=\\,05^h01^m57^s$, $\\delta_{2000}\\,=\\,11^\\circ46'24''$. Eight hours after the GRB trigger, from February 28.4681 to February 28.8330 UTC, the Narrow Field Instruments (NFI) on board \\sax\\ (Boella et al. 1997a) were pointed to the WFC error box. An X-ray source, SAX J0501.7+1146, was detected (Costa et al 1997b) in the field of view of both the Low Energy (0.1--10~keV) and Medium Energy (2--10~keV) Concentrators Spectrometers (LECS and MECS) (Parmar et al. 1997; Boella et al. 1997b). The source position ($\\alpha_{2000}\\,=\\,05^h01^m44^s$, $\\delta_{2000}\\,=\\,11^\\circ46'42''$) is consistent with the GRB error circle. The source was again observed about three days later, from March 3.7345 to March 4.1174. During this observation, the 2--10~keV source flux had decreased by about a factor 20, while in the 0.1--2~keV energy range the source was not detected. Following the discovery of the GRB, searches for radio and optical counterparts to GRB970228 were conducted with most of the ground based telescopes in the northern hemisphere. Groot et al. (1997) reported the discovery of an optical transient at a position ($\\alpha_{2000}\\,=\\,05^h01^m46.70^s$, $\\delta_{2000}\\,=\\,11^\\circ46'53.0''$), consistent with both the \\sax\\ WFC and NFI error boxes and with the long baseline timing \\ulysses/\\sax\\ and  \\ulysses/\\wind\\ error annuli, of 31$''$ and 30$''$ half-width, respectively (Hurley et al. 1997; Cline et al. 1997). While the association of the transient X-ray source with the afterglow of GRB970228 was compelling on the basis of the properties of its decay curve when extrapolated backwards to the burst time (Costa et al. 1997c), the association of the optical transient with the burst afterglow was less strong. In spite of the positional consistency and temporal behaviour of the optical transient, it was not possible to exclude the possibility that the optical transient was unrelated to the GRB (see discussion by van Paradijs et al. 1997), like in the case of the radio source discovered in the earliest error box of GRB970111, which showed  a time behaviour consistent with that expected from radio afterglows  of GRBs, but later resulted to be unrelated tho the burst  (Frail et al. 1997). The \\rosat\\ satellite, thanks to its HRI focal plane instrument, offered the possibility of imaging the X-ray afterglow at 10$''$ angular resolution (David et al. 1997). A Target of Opportunity observation was thus requested and obtained. Here we report on results of that observation and its consequences. Preliminary results were already previously reported (Frontera et al. 1997b). ",
        "conclusions": "The \\rosat\\ HRI observation of the \\sax\\ WFC error circle of GRB970228 clearly shows the presence of a new X-ray source. Its position  within the error box of the \\sax\\ source, the low probability of a chance coincidence ($\\sim 1 \\times 10^{-3}$) and the better imaging capabilities of the \\rosat\\ HRI compared to the \\sax\\ NFI, indicate that the \\rosat\\ source and the \\sax\\ source are the same object. The source position derived from the \\rosat\\ observation is the most precise position of a GRB X-ray afterglow obtained thus far. Its position is also coincident with the optical transient associated to GRB970228 within 2$''$. This result confirms that X-ray source and the optical transient are the same object. The X-ray source hows a 0.1--2.4 keV decline according to a power law decline with index $\\alpha \\,=\\, 1.50^{+0.23}_{-0.35}$ for at least 13 days. This slope is fully consistent with that estimated in the 2--10 keV energy band (Costa et al. 1997c) and is marginally consistent with that reported by Fruchter et al. (1997, 1998) in the optical band. Thus it appears that from the X-ray to the optical band the GRB afterglow has the same decline law. \\begin{figure} \\epsfig{figure=grb282_l_tot.ps,height=8.5cm,width=8.5cm,angle=0} \\caption{As in Fig. 2, but extrapolated to the first second from the burst onset. The two arrows on the left delimit the time interval, that corresponds to the GRB last three pulses, when the X-ray afterglow is expected to start (see text). } \\label{figure:decayb} \\end{figure}"
    },
    "9804/astro-ph9804264_arXiv.txt": {
        "abstract": "Including nucleon--nucleon correlations due to both Fermi statistics and nuclear forces, we have developed a general formalism for calculating the charged--current neutrino--nucleon absorption rates in nuclear matter.  We find that at one half nuclear density many--body effects alone suppress the rates by a factor of two and that the suppression factors increase to $\\sim$5 at $4\\times10^{14}$ g cm$^{-3}$.  The associated increase in the neutrino--matter mean--free--paths parallels that found for neutral--current interactions and opens up interesting possibilities in the context of the delayed supernova mechanism and protoneutron star cooling. ",
        "introduction": "The neutrino absorption and scattering opacities in the post--shock core of a supernova, in which nuclei are largely disintegrated into nucleons, determine the duration, spectrum, and flavor distribution of the emerging neutrino pulse. It has been known for some time that the interactions among nucleons in the denser regions can change these opacities significantly, but to date there has been no comprehensive treatment given in the literature and present calculations of the complete supernova process do not include the effects of interactions on the opacities.  The neutrino--matter interaction rates can be related to the space-- and time--dependent correlations among the set of density operators for the separate nuclear constituents (to find the Gamow--Teller parts we must consider separate spin--up and spin--down densities). In the case of neutral--current interactions \\cite{BS}, there is an instructive limit, which also provides an estimate of the effects, in which the combined limits of large nucleon mass and small neutrino energy allow the use of long--wavelength limits of equal--time correlation functions, in turn expressible in terms of the second derivatives of an energy density functional with respect to various densities. This approach is the direct multichannel generalization of the familiar results for light scattering from the thermal density fluctuations in a fluid, where it is the compressibility that determines the long--wavelength opacity, and it was used in references \\cite{s2Ray} and \\cite{iwa} to find significant reductions of neutral--current opacity in certain regions. In Burrows \\& Sawyer \\cite{BS}, an approach based on ring graphs was used to encompass these results and to extend them to domains in which the equal--time and long--wavelength limits are not clearly applicable.  The use of the equal--time and long--wavelength limits to express correlation functions in terms of static susceptibilities cannot be extended to the charged--current interactions when there is a large chemical potential difference between protons and neutrons. Furthermore, there do not exist in the present literature systematic estimates of the effects of interactions on the charged--current opacities for electron neutrinos. In the present work, we give a theoretical framework for addressing these opacities, based on summing ring graphs, together with the results of calculations with input parameters taken from the current phenomenology of nuclear matter. ",
        "conclusions": "We have developed a new formalism for incorporating the effects of many--body correlations on the charged--current rates of neutrino--matter interactions. This formalism reveals that these rates are considerably suppressed in the densest regions of protoneutron stars and supernova cores.  Assuming that the nucleons are non--relativistic, our formalism incorporates the full kinematics of the interaction, Pauli blocking by final--state nucleons (protons), and correlation due to nucleon--nucleon interactions.  We have employed the ring approximation (RPA) and assumed the near--validity of Fermi Liquid Theory.  It would desirable to include ladder diagrams and to perform the calculations in the context of a better numerical method for solution of the nuclear equation of state (EOS), since the solution of the EOS is intimately related to the derivation of the scattering/absorption rates. However, those who perform detailed nuclear EOS calculations and address many--body correlations in nuclear matter do not as yet provide the requisite spin and density structure functions, even for the static case.  These results for charged currents, when combined with the results from Burrows \\& Sawyer \\cite{BS} for neutral currents, strongly suggest that energy and lepton number will leak from supernova cores at a rate that is higher than heretofore estimated. This implies that the neutrino luminosities during the epoch after bounce for which the inner core is the major energy source ($> 0.5 - 1.5$s) will be enhanced, perhaps by as much as 50\\% \\cite{BS}.  The consequences of this increased transparency for the neutrino--driven supernova explosion mechanism \\cite{bhf} may be interesting, but have yet to be clarified."
    },
    "9804/astro-ph9804052_arXiv.txt": {
        "abstract": "ACO 3627 is a rich, nearby cluster of galaxies at the core of the Great Attractor. At the low galactic latitude of $b = -7.2\\deg$ the galactic extinction is significant. Nevertheless, its proximity makes it a prime target for studies of environmental effects on its cluster members. Here, we report on a multi-wavelength study of a Seyfert 1 galaxy at 30 arcmin from the centre of ACO 3627. Its Seyfert nature was discovered spectroscopically and confirmed in X-rays. We have obtained B$_{\\rm J}$ and  R$_{\\rm C}$ CCD photometry as well as J, H, K and L aperture photometry at the SAAO, low and high resolution spectroscopy (ESO and SAAO), 21 cm line observations (Parkes Observatory) and X-ray ROSAT PSPC data.  The Seyfert 1 galaxy is of morphology SBa(r). It has a nearby companion (dS0) but shows no signs of interaction. A consistent value for the galactic extinction of A$_{\\rm B}$ = 1.6 mag could be determined. The nucleus of the Seyfert is very blue with a strong (B$_{\\rm J}$ -- R$_{\\rm C}$) colour gradient in the inner 2.5 arcsec. The extinction-corrected near-infrared colours of WKK 6092 are typical of a Seyfert 1 and the X-ray spectrum conforms to the expectation of a Seyfert as well. The galaxy has a very low \\HI\\ flux. This could be explained by its morphology, but also -- due to its very central position within the rich Norma cluster -- to ram pressure stripping. ",
        "introduction": "Dust and stars in the Milky Way obscure a large fraction of the extragalactic sky, creating a ``Zone of Avoidance'' (ZOA) in the distribution of galaxies. In an effort to reduce the size of the ZOA and thus coming closer towards an all-sky distribution of galaxies, we have embarked on a deep optical galaxy search behind the southern Milky Way (Kraan-Korteweg \\& Woudt 1994). This has led to the recognition that ACO 3627 (Abell \\etal\\ 1989), also called the Norma cluster after the constellation it is located in, is a massive, nearby cluster of galaxies at the core of the Great Attractor (GA) $(\\ell,b,v) = (325\\deg, -7\\deg, 4882$ \\kms) (Kraan-Korteweg \\etal\\ 1996). The Norma cluster appears to be the central, dominant component of a ``great wall''-like structure and would be the most prominent overdensity of galaxies in the southern sky, were it not obscured by the Milky Way (Woudt \\etal\\ 1997).  Recent observations of the Norma cluster with the ROSAT PSPC have confirmed the massive nature of ACO 3627; it is the 6$^{th}$ brightest cluster in the ROSAT sky (B\\\"ohringer \\etal\\ 1996). The X-ray contours furthermore suggest the existence of a subcluster. The merging scenario is independently supported by the radio continuum emission of the central cD galaxy PKS1610-608. The emission from this wide-angle-tail (WAT) radio source  (Jones \\& McAdam 1992) seems to encircle the X-ray subcluster (\\cf\\ Fig.~3 of Kraan-Korteweg \\etal\\ 1997) and is indicative of a strong motion of the cluster gas due to the ongoing merging process (Jones \\& McAdam 1996, Burns \\etal\\ 1994).  Roughly 30{\\arcmin} from the centre of this cluster -- taken as the central cD galaxy PKS1610-608 -- we have identified a Seyfert 1 galaxy. It is a member of ACO 3627. In the following sections we describe the various observations of this galaxy: the discovery of the galaxy in section 2, the multicolour photometry obtained at the South African Astronomical Observatory (SAAO) in section 3, the spectroscopy obtained at the European Southern Observatories (ESO) and the SAAO in section 4, the \\HI\\ observations obtained with the 64m radio telescope of the Parkes Observatory of the Australian Telescope National Facility (ATNF) in section 5, and the X-ray data from ROSAT PSPC observations in section 6. The results are summarized and discussed in the last section. ",
        "conclusions": "We have observed the Seyfert galaxy WKK 6092 at different wavelengths. The resulting data are summarized in Table~\\ref{seytab}.  \\begin{table} \\caption{Observational parameters of WKK 6092} \\label{seytab} \\begin{tabbing} Apparent blue (IIIaJ)(m$_{B}$) magn  \\=    \\kill  {\\bf Coordinates:} \\> \\\\ R.A. (1950)        \\> $16^{h} 07^{m} 32.7^{s}$ \\\\ DEC. (1950)        \\> $-60^{\\circ} 30' 11''$ \\\\ Galactic Longitude \\> $325.20^{\\circ}$ \\\\ Galactic Latitude  \\> $-6.74^{\\circ}$ \\\\  \\vspace{5mm} {\\bf Properties:} \\> \\\\ Hubble Type        \\> SBa(r) \\\\ Dimensions (a x b) \\> 56'' x 47'' \\\\ Ellipticity (1-b/a)  \\> 0.11 \\\\ Inclination \\> 28$\\deg$ \\\\ Position Angle \\> 96$\\deg$  \\\\  \\vspace{5mm} {\\bf Photometry:} \\> \\\\ $B_{\\rm J}$ (IIIaJ)      \\> 14.7 $\\pm$ 0.5 mag \\\\ $B_{25}$ (CCD)     \\> 14.96 $\\pm$ 0.09 mag \\\\ $B_{\\rm T}$ (CCD)      \\> 14.88 $\\pm$ 0.13 mag \\\\ $R_{24}$ (CCD)     \\> 13.38 $\\pm$ 0.12 mag \\\\ $R_{\\rm T}$ (CCD)      \\> 13.30 $\\pm$ 0.14 mag \\\\ J$_{\\rm c}$              \\> 12.91 $\\pm$ 0.03 mag \\\\ H                  \\> 11.98 $\\pm$ 0.03 mag \\\\ K                  \\> 11.51 $\\pm$ 0.03 mag \\\\ L                  \\> 10.60 $\\pm$ 0.20 mag \\\\ \\HI\\ flux          \\> 0.93 Jy \\kms \\\\ X-Ray (0.5--2.0 keV): \\>    \\\\ \\hspace{0.25cm} Flux \\> $1.05 \\pm 0.15 \\cdot 10^{-12}$ erg s$^{-1}$ \\\\ \\>   \\hspace{2.65cm}       cm$^{-2}$ \\\\ \\hspace{0.25cm} ${\\cal L}_X$ \\> $1.2 \\pm 0.17 \\cdot 10^{42}$ erg s$^{-1}$\\\\  \\vspace{5mm} {\\bf  Galactic Extinction (A$_B$):} \\>  \\\\ from HI      \\> 1.5 mag \\\\ from Balmer decrement \\> $\\la$ 1.7 mag \\\\ from X-ray   \\> 1.6 mag \\\\  \\vspace{5mm} {\\bf Heliocentric velocity:}       \\> \\\\ MEFOS        \\> 4711 $\\pm$ 30  \\kms \\\\ S.A.A.O.     \\> 4688 $\\pm$ 40 \\kms \\\\ Parkes 64-m  \\> 5012 $\\pm$ 5 \\kms \\\\ \\hspace{0.25cm} \\DVF         \\> 88 \\kms \\\\ \\hspace{0.25cm} \\DVT         \\> 97 \\kms \\\\  \\end{tabbing} \\end{table}  WKK 6092 and its neighbour are both members of ACO 3627. They have similar redshifts but show no indications of interaction. The morphology of both galaxies do not seem distorted (\\cf\\ Fig.~\\ref{ccdim}). The Seyfert has a very blue nucleus, a distinct bar and a ring superimposed on an otherwise smooth disk.  An upper limit for the Galactic foreground extinction in the line of sight of the Seyfert galaxy can be set at A$_{\\rm B}$ = 1.6 mag. This was determined by three different methods, the Balmer decrement in the optical spectrum, the fitting of an absorbed power low to the X-ray spectrum and the Galactic \\HI\\ column densities. All give a consistent value of the foreground extinction. A minor fraction of the extinction is intrinsic to the galaxy.  The extinction corrected near-infrared colours of WKK 6092 are typical of a Seyfert 1 and are in agreement with well known Seyfert 1's such as NGC 1566 (Glass and Moorwood 1985). The X-ray sepctrum is also consistent with the standard expectation for this object.  At the adopted cluster distance of R = 93 $h_{50}^{-1}$ Mpc, the absolute magnitude (corrected for the galactic extincton) is M$_{B_T}^o = -21.52$. The \\HI\\ and total mass is $1.9 \\cdot 10^{9} {\\cal M}_{\\odot}$ and $30 \\cdot 10^{9} {\\cal M}_{\\odot}$, respectively. The Seyfert is at a projected distance of 0.8 $h_{50}^{-1}$ Mpc from the cluster centre and the \\HI\\ content of the galaxy might be influenced by interactions with the Inter Cluster Medium due to processes like ram pressure stripping (Cayatte \\etal\\ 1990). The \\HI\\ content is in fact quite low. This is, however, not inconsistent with the expectation for a barred early-type spiral.  Despite the difficulties in analysing data of an object deep within the Milky Way, all data concerning the here investigated Seyfert galaxy WKK 6092 at 30 arcmin from the centre of the rich cluster ACO 3627 correspond to the standard characteristics of a Seyfert 1 galaxy."
    },
    "9804/astro-ph9804322_arXiv.txt": {
        "abstract": "We report the {\\em first} unambiguous detection of the host galaxy of a normal radio-quiet QSO at high-redshift in $K$-band. The luminosity of the host comprises about 35\\%\\ of the total $K$-band luminosity. Assuming the average colour of QSOs at $z\\approx2$, the host would be about 5 to 6~mag brighter than an unevolved $L_*$ galaxy placed at $z\\approx2$, and 3 to 4~mag brighter than a passively evolved $L_*$ galaxy at the same redshift. The luminosity of the host galaxy of the QSO  would thus overlap with the highest found in radio-loud QSOs and radio-galaxies at the same redshift. ",
        "introduction": "Recent evidence that the cosmological evolution of the density of star formation in the Universe  (Madau et al. 1996) follows closely the QSO density evolution (Boyle \\& Terlevich 1998) emphasizes the need to study the kinds of galaxies that host Active Galactic Nuclei in order to understand the link between star-formation and nuclear activity, and potentially the role of nuclear activity in galaxy formation.  At the peak value of QSO density ($z\\approx 2$ to 3), the few QSO host-galaxies detected so far present rest-frame UV fluxes that reach up to 20\\%\\ of the total QSO luminosity, indicating star-formation rates about 200 \\Msun/yr and above for both radio-loud (Lehnert et al. 1992) and radio-quiet samples (Aretxaga, Boyle \\& Terlevich 1995, Hutchings 1995). These values are almost an order of magnitude above those of field galaxies at similar redshifts selected through Lyman Break techniques (Steidel et al. 1996, Lowenthal et al. 1997). The properties of these QSO hosts are not unprecedented, since they follow very closely the luminosity--size relation of nearby star forming galaxies, overlapping with its high-luminosity end (Aretxaga, Terlevich \\& Boyle 1998). However, the UV fluxes only carry information about the high-mass end of the stellar populations in the galaxies, and say little about the bulk of the stellar mass which is better characterized by optical to NIR observations.  Although a few hosts of extreme radio-loud QSOs at $z\\approx 2$ have been detected in NIR bands (Lehnert et al. 1992, Carballo et al. 1998), attempts to image the hosts of normal radio-quiet QSOs at the same redshifts have been unsuccessful to date (Lowenthal et al. 1995, Aretxaga et al. 1998). Imaging radio-quiet systems, which constitute more than 95\\%\\ of all QSOs, is important in order to characterize the bulk of the population. The observed optical sizes of FWHM$\\approx 1$ arcsec (Aretxaga et al. 1995), clearly demand a technique which offers the highest available angular resolution.  In this paper we focus our attention on the detection of the host of a normal radio-quiet $z\\approx 2$ QSO with the Adaptive Optics System in operation at the ESO 3.6m telescope in La Silla. Preliminary results on similar programs to image the host-galaxies of QSOs at $z\\approx 0.5 \\hbox{ \\ and \\ } 1.7$ using Adaptive Optics have been presented in a recent conference devoted to quasar hosts (Bremer et al. 1997, Hutchings 1997).     \\ifoldfss ",
        "conclusions": ""
    },
    "9804/astro-ph9804114_arXiv.txt": {
        "abstract": "We obtain self-similar solutions that describe the gravitational collapse of nonrotating, isothermal, magnetic molecular cloud cores.  We use simplifying assumptions but explicitly include the induction equation, and the semianalytic solutions we derive are the first to account for the effects of ambipolar diffusion following the formation of a central point mass. Our results demonstrate that, after the protostar first forms, ambipolar diffusion causes the magnetic flux to decouple in a growing region around the center. The decoupled field lines remain approximately stationary and drive a hydromagnetic C-shock that moves outward at a fraction of the speed of sound (typically a few tenths of a kilometer per second), reaching a distance of a few thousand AU at the end of the main accretion phase for a solar-mass star.  We also show that, in the absence of field diffusivity, a contracting core will not give rise to a shock if, as is likely to be the case, the inflow speed near the origin is nonzero at the time of point-mass formation. Although the evolution of realistic molecular cloud cores will not be exactly self similar, our results reproduce the main qualitative features found in detailed core-collapse simulations (Ciolek \\& K\\\"{o}nigl 1998). ",
        "introduction": "Low-mass stars are generally believed to form as a result of the gravitational collapse of molecular cloud cores. The cores are initially supported by thermal and magnetic forces, but because of ambipolar diffusion (the drift of ions, to which the magnetic field lines are attached, relative to the dominant neutral gas component), they gradually lose their magnetic support and eventually collapse after becoming ``supercritical'' (see, e.g., Mouschovias 1987 for a review).\\footnote{In this paper we reserve the term ``core'' for the high-density central region of a molecular cloud and do {\\em not} apply it to the point mass that forms from the collapse of such a core.} The most detailed numerical treatments to date of the problem of the ambipolar diffusion-initiated formation of supercritical cores and the early stages (prior to point mass formation) of their subsequent dynamical collapse have been presented by Mouschovias and collaborators (Fiedler \\& Mouschovias~1992, 1993; Ciolek \\& Mouschovias~1993, 1994, 1995, hereafter CM93, CM94, CM95; Basu \\& Mouschovias~1994, 1995a, 1995b, hereafter BM94, BM95a,b). Because the timescale for core formation is much longer than the timescale for dynamical collapse, special numerical techniques had to be employed in these calculations. The simulations were terminated when the central densities reached $\\sim 10^{10}\\ {\\rm cm}^{-3}$ and the underlying assumptions of isothermality (e.g., Gaustad~1963) and flux freezing onto the ions (e.g., Pneuman \\& Mitchell~1965) broke down.  These calculations were nevertheless able to demonstrate that {\\em supercritical cores begin to collapse dynamically before a point mass (i.e., a protostar) appears at the origin}.  The dynamical evolution of supercritical cores after their formation has been studied by many researchers. Solutions exist for the collapse of nonrotating, self-gravitating spheres without thermal support (Henriksen~1994), self-gravitating spheres with thermal support (Penston~1969; Larson~1969; Shu~1977; Hunter~1977; Boss \\& Black~1982; Whitworth \\& Summers~1985; Foster \\& Chevalier~1993) as well as with a combined thermal and isotropic magnetic pressure support (Chiueh \\& Chou 1994), and self-gravitating disks with thermal support (Narita, Hayashi, \\& Miyama~1984; Matsumoto, Hanawa, \\& Nakamura~1997) and also with ordered, frozen-in magnetic fields (Nakamura, Hanawa \\& Nakano~1995; Li \\& Shu~1997, hereafter LS). In order to choose a particular solution for a given problem, one needs to know the properties of the supercritical core at the time of its formation. This information, however, can only be gleaned from a study of the preceding, quasi-static evolution of the core under the influence of ambipolar diffusion. Although different assumptions about the initial state of the core yield solutions that are qualitatively similar in their gross behavior (the core collapses with near free-fall speeds and a point mass eventualy forms at the center), the solutions do differ in such important details as the accretion rate onto the central point mass and the formation (or absence) of shocks.  The well-known examples of the Larson-Penston (1969) and Shu~(1977) similarity solutions in fact represent two extremes of a whole continuum of self-similar collapse solutions specified by a cloud's initial configuration and the conditions at its boundary (Hunter 1977; Whitworth \\& Summers 1985; see also Chiueh \\& Chou 1994 for a generalization to the case of an isotropic internal magnetic pressure). The Larson-Penston (1969) solution is characterized by a spatially uniform, supersonic (at $\\sim 3.3$ times the isothermal speed of sound $C$) infall speed and an inverse-square dependence of the density $\\rho$ on the radius $r$ at the instant of point-mass formation (PMF); the mass accretion rate at the center is $\\sim 29 \\ C^3/G$ (where $G$ is the gravitational constant) at that instant and increases to $\\sim 47 \\ C^3/G$ immediately after PMF. Numerical simulations of the collapse of nonmagnetic isothermal spheres (Hunter 1977; Foster \\& Chevalier 1993) have indicated that this solution provides a good approximation to the conditions near the center at the PMF epoch for clouds that are initially near a marginally stable equilibrium. The Shu (1977) solution strictly applies only to the post-PMF evolutionary phase: it consists of an inner free-fall region and a hydrostatic outer envelope that are separated by an outward-propagating (at a speed $C$ relative to the gas) expansion wave. The envelope corresponds to a singular isothermal sphere ($\\rho \\propto r^{-2}$) and the mass accretion rate onto the center is $\\sim 1 \\ C^3/G$. In applying this solution to real systems, it was proposed to identify the initial core configuration at the end of the quasi-static ambipolar-diffusion phase with a singular isothermal sphere (or, more generally, a toroid) at the instant of PMF (e.g., Shu, Adams, \\& Lizano 1987; Li \\& Shu 1996). However, as we noted above, the conclusion from detailed numerical simulations has been that the dynamical phase of core collapse generally commences well before the PMF epoch, so that the innermost region is not well represented by a quasi-static solution at the time of point-mass formation.  Another interesting effect that depends on the specific choice of initial conditions and on the detailed physical properties of the collapsing core is the formation (or absence) of shocks (e.g., Tsai \\& Hsu 1995; LS).  For example, LS discovered that when, instead of a spherical core, one considers the collapse of a flattened disk, the expansion wave of Shu~(1977) becomes a shock. As we show in this paper, when one takes proper account of the fact that supercritical cores collapse dynamically before a point mass first forms at the origin, that shock disappears.  Nevertheless, a physical basis for the formation of shocks in collapsing magnetized molecular cloud cores has been discussed by Li \\& McKee (1996), who argued that a hydromagnetic C-shock will appear as a result of the outward diffusion of inwardly advected magnetic flux. The existence of such a shock has been confirmed in the numerical simulations of Ciolek \\& K\\\"{o}nigl (1998, hereafter CK), and it is, in fact, a salient feature of the semianalytic solutions derived in this paper.  The aim of the present work is to clarify the effects of ambipolar diffusion in dynamically collapsing supercritical cores. Toward this goal, we construct semianalytic, time-dependent similarity solutions of gravitationally contracting, magnetized, isothermal disks. Although the evolution of real molecular cloud cores is not expected to be exactly self similar, we demonstrate, through a comparison with the detailed numerical simulations of CK, that our solutions capture the main traits exhibited by the latter calculations. Based on an analogous comparison with the results of numerical simulations, Basu (1997) showed that a self-similar scaling describes the pre-PMF evolution in the innermost flux tubes of collapsing supercritical cores quite well. To complement his study, we concentrate in this paper on the post-PMF evolutionary phase. Our approach differs, however, from that of Basu (1997) in that we explicitly solve the induction equation, whereas he accounted for the effects of ambipolar diffusion only in a phenomenological manner.\\footnote{Our work is thus also distinguished from that of Safier, McKee, \\& Stahler (1997), who studied the effects of ambipolar diffusion in the spherically symmetric, quasi-static limit without explicitly solving the induction equation.} In fact, the solutions that we derive, while involving various simplifications, are nevertheless the first to consistently incorporate ambipolar diffusion into a self-similar representation of the collapse of a magnetized cloud core. \\footnote{The effect of {\\em weak} magnetic fields on a dynamically collapsing core in the presence of ambipolar diffusion was previously investigated by Galli \\& Shu (1993a), who carried out a perturbation expansion of the (nonmagnetic) spherical similarity solution of Shu (1977). As was already noted and discussed by Li \\& McKee (1996), the semianalytic solution derived in that paper, as well as the associated numerical calculation in Galli \\& Shu (1993b), did not uncover the existence of a flux diffusion-driven shock.} We formulate the problem in \\S 2, present our solutions in \\S 3, and discuss the results in \\S 4. Our conclusions are summarized in \\S 5. ",
        "conclusions": "In this paper we have presented a self-similar solution of the collapse of a magnetized molecular cloud core (assumed to also be nonrotating and isothermal) that, for the first time, incorporated the effects of ambipolar diffusion in a self-consistent manner. We have focused on the post-PMF (point-mass formation) phase of the collapse of a disk-like core, noting that Basu (1997) had previously explored the self-similar nature of the collapse before a central mass (i.e., a protostar) first appears at the origin. We clarified the distinction between the ideal and nonideal MHD cases by plotting the singular lines in the position--velocity space and showing that they correspond to different critical speeds (the magnetosonic speed and thermal sound speed in the ideal and nonideal problems, respectively). We obtained a solution for the ideal (flux-frozen) case that exhibits a split-monopole field topology near the center. This solution differs from the one obtained by Li \\& Shu (1997) in that it involves no shocks. We showed that the shock in the LS solution is a direct consequence of their assumption that the core at the time of PMF is described by a stationary density distribution (corresponding to a singular isothermal toroid), and we pointed out that a shock will generally {\\em not} be present under the more realistic assumption of a nonzero inflow speed near the origin at that instant. We demonstrated, however, that a shock is a generic feature of the solution in the nonideal (ambipolar diffusion) case. This (C-type) shock is a direct consequence of the action of ambipolar diffusion in the central region of the core following PMF: the magnetic diffusivity decouples the field from the matter, causing the gas to free-fall to the center (where it accumulates in a point mass) and the field to stay behind and drive a shock outward. We have compared this solution with the results of the numerical simulations of Ciolek \\& K\\\"onigl (1998) and confirmed that, while the more realistic numerical models are not strictly self-similar, our simplified solution nevertheless captures the main features of the core evolution after PMF."
    },
    "9804/astro-ph9804138_arXiv.txt": {
        "abstract": "s{ We review the topic of Cosmic Microwave Background Anisotropy measurements carried out by means of balloon-borne telescopes. After a short description of the experimental methodology, we outline the peculiar problems of these experiments, and we describe the main results obtained and the perspects for future developments. } ",
        "introduction": " ",
        "conclusions": "Ballooning for CMB Anisotropy measurements is a very active field worlwide. The activity is growing, for two main scientific reasons:  1) High frequency ($>$90 GHz) and high angular resolution ($\\sim$ 10 arcmin FWHM) measurements are possible and effective.  2) These measurements complement the forthcoming data from MAP and are a very important test-bed for Planck technologies.  Moreover, good science is being produced, and promises for important results (like the $\\ell$-space spectroscopy of the acoustic peaks, or detection/falsification of non-gaussian statistics for the CMB fluctuations) are quite convincing.  As an example relevant for this conference, we can mention the fact that determination of several cosmological parameters is possible with very good accuracy from LDB experiments. For example, if all the systematics effects are properly removed, a single LDB experiment with 12 arcmin FWHM beams, 16 total power bolometric detectors with sensitivity of 80 $\\mu K \\sqrt{s}$, 10 days of observing time spent over a $50^o \\times 50^o$ sky region, can measure the power spectrum of the CMB anisotropies with very good accouracy \\cite{Silvia}. Fits can be done on the measured power spectrum \\cite{KL}, allowing to recover $\\Omega_{tot}$ with a 3$\\%$ error, $\\Omega_{\\Lambda}$ with 6$\\%$ error, $\\Omega_B$ with $1\\%$ error, $n$ scalar with 18$\\%$ error, $H_o$ with 1$\\%$ error, $Q_{rms,PS}$ with 4$\\%$ error. Here the errors for any parameter make no assumptions about the value of the other parameters. These measurement errors can be significantly reduced if one or more of the cosmological parameters are constrained by other observations or fixed by assumptions."
    },
    "9804/astro-ph9804281_arXiv.txt": {
        "abstract": "We present a description of the observations and data reduction procedures for an extensive spectroscopic and multi-band photometric study of nine high redshift, optically-selected cluster candidates. The primary goal of the survey is to establish new constraints on cluster and galaxy evolution, with specific emphasis on the evolution of galaxy morphology and on the star-formation history of the galaxies within and around distant clusters. We have measured 892 new redshifts for galaxies with $R \\le 23.3$. The data will also serve as deep probes of the foreground and background large-scale structures. The observations include broad band optical imaging and spectroscopy with the Low Resolution Imaging Spectrograph at the 10 meter W. M. Keck Observatory telescope; K-band imaging with IRIM at the 4 meter Kitt Peak National Observatory telescope; and deep, high angular resolution imaging with the WFPC2 onboard the Hubble Space Telescope. We also describe the procedures used to obtain morphological information. We have established that six of the nine cluster candidates are indeed real space density enhancements and are representative of those typically associated with clusters of galaxies. The remaining three candidates appear to be projections of several smaller groups at widely separated distances. This success rate is consistent with estimates of the false positive rate in 2D optical high-$z$ cluster searches. ",
        "introduction": "Clusters of galaxies have historically provided an important tool for studying cosmology and the evolution of galaxies. Because of their high concentration of galaxies, clusters provide an environment in which to study large, statistical samples of galaxies. Therefore, examining clusters of galaxies from the local universe to those at high redshift allows us to probe galactic evolution to redshifts of the order of 1. Clusters of galaxies at redshifts of $z \\simless 0.2$ have been extremely well cataloged in the optical regime (e.g.\\ Abell 1958; Zwicky \\etal 1968; Dressler 1980; Shectman 1985; Abell \\etal 1989; Lumsden \\etal 1992; Dalton \\etal 1994). The analyses of local clusters indicate that these systems are dense (Abell richnesses of $\\sim 30 - 300$ galaxies), massive ($M \\sim 10^{14} - 2 \\times 10^{15}~h^{-1}~{\\rm M_{\\odot}}$), and dominated by early-type galaxies ($\\sim 50 - 80\\%$ of the total galaxy population). These studies provide a strong basis on which to compare cluster properties at increasingly higher redshift.  Detailed photometric, spectroscopic and morphological studies have been extended to clusters of galaxies at redshifts up to $z \\sim 0.6$. The first substantial contribution at these redshifts came from Butcher \\& Oemler (1984) who found a surprisingly large population of blue galaxies in conjunction with the expected red sequence of early-type cluster members. Further photometric and spectroscopic campaigns, including the ambitious CNOC and MORPHS surveys, have confirmed the progressive bluing of the cluster's galaxy population and have tracked the passive evolution of the early-type galaxies (Dressler \\& Gunn 1992; Oke, Gunn \\& Hoessel 1996; Yee, Ellingson \\& Carlberg 1996; Ellingson \\etal 1997; Ellis \\etal 1997; Stanford et al.\\ 1995, 1997; and references therein). The Hubble Space Telescope (HST) has enabled morphological classification of intermediate-redshift ($z \\simless 1$) galaxies on scales which are comparable to the classifications made of their local counterparts. These high-resolution studies have revealed that there may be a substantial change in the morphological composition of the clusters (Smail \\etal 1997; Dressler \\etal 1997; however, see Stanford \\etal 1997). All of these results imply that there is a significant amount of evolution occurring in the cluster environment between redshifts of $z \\approx 0.5$ and $z = 0.0$.  In order to understand this apparent change in the galaxy population, it is essential to probe in similar detail clusters of galaxies at even higher redshift where the effects of evolution and cosmology will be even greater.  In light of this, we have undertaken an extensive survey of nine candidate clusters of galaxies at redshifts greater than 0.6. Only a few optical/near-IR surveys have attempted to detect systematically clusters at high redshift; therefore, we have chosen our cluster sample from the Gunn, Hoessel, \\& Oke (1984) survey and the Palomar Distant Cluster Survey (Postman \\etal\\ 1996).  The observational data compiled for this survey, to date, includes deep $BVRIK$ photometry, over 900 low-resolution spectra, and deep F606W/F702W/F814W imagery from HST. The large redshift database allows us to reliably distinguish between physically real clusters and chance line-of-sight projections. The full data allow us to measure the global properties of the clusters, such as profile shape and dynamics, as well as the individual properties of the cluster galaxies, such as color, star-formation rate, and morphology.  In this introductory paper to our high-redshift cluster series, we describe in detail the observational and data reduction techniques of each aspect of this survey and present the redshift histograms for our nine fields. The subsequent papers in this series will present the specific analyses and scientific results of this survey. These papers include the second and third installments in this series which present a detailed photometric, spectroscopic, and morphological analyses of the first two clusters to be completed in this survey, CL0023+0423 and CL1604+4304 (Postman, Lubin \\& Oke 1998; Lubin \\etal 1998). ",
        "conclusions": "We have described the data acquisition and reduction procedures of our photometric and spectroscopic campaign to study nine candidate clusters of galaxies at redshifts of $z > 0.6$. The observational program consists of four main parts :  \\newcounter{discnt} \\begin{list} {\\arabic{discnt}.} {\\usecounter{discnt} \\setlength{\\leftmargin 0.2in}{\\itemsep 0in}{\\topsep 0in}{\\parskip 0in}}  \\item Spectra for approximately 80\\% of all galaxies down to a Johnson--Cousins $R$ magnitude of $\\sim 23.5$ within a fixed area of ${2}^{'}.2 \\times {7}^{'}.6$ of each cluster field using the Low--Resolution Imaging Spectrograph (LRIS) at the Keck 10m telescope. We have obtained spectra covering the range 4400 \\AA\\ to 9500 \\AA\\ for $\\sim 130 - 150$ galaxies per cluster field. Redshifts have been determined for 892 objects.  \\item Deep $BVRI$ imaging with Keck of all galaxies in the full LRIS field of $6^{'} \\times 8^{'}$ centered on each cluster.  The $5\\sigma$ detection limits are $B = 25.1$, $V = 24.1$, $R = 23.5$, and $I = 21.7$ in our standard 3 arcsecond radius aperture.  The photometric data are converted to absolute fluxes in order to obtain absolute spectral energy distributions.  \\item High angular resolution imagery with HST in order to provide morphological information on the galaxies in the WFPC2 field-of-view (${160}^{''} \\times {160}^{''}$) centered on each cluster. Each of the cluster candidates has been or will be observed by HST in Cycles 5 and 6.  \\item High precision $K$ band photometry across the WFPC2 field--of--view with the KPNO 4m telescope for each cluster.  The infrared survey reaches a limiting magnitude of $K^{'} = 20$ in the standard aperture. \\end{list}  We have presented the redshift histograms for the nine candidate clusters of galaxies in this survey. We find that six of the nine candidate clusters of galaxies are real density enhancements.  They include CL0023+0423, CL0943+4804, CL1324+3011, CL1325+3009, CL1604+4304, and CL1604+4321. The remaining three candidates are of a more dubious nature. Their redshift distributions reveal no clear density enhancement but rather an apparent superposition of small groups of galaxies along the line-of-sight. This false positive rate is consistent with the estimate of $\\sim 30$\\% provided in Postman \\etal (1996). At lower redshifts ($z \\simless 0.5$) the spurious rate is about 20\\% or less. This is based on spectroscopic follow-up of PDCS candidates being done by Holden \\& Nichol (1998) at the KPNO 4m telescope. We conclude that optical detection of clusters remains a successful and important method for identifying such systems out to $z \\sim 1$ and, further, will provide an important complement to cluster searches at other wavelengths.  Results on the star-formation history, dynamics, and morphological properties of  CL0023+0423 ($z = 0.84$) and CL1604+4304 ($z = 0.90$) are presented in Papers II and III.  \\vskip 1.0cm We thank Don Schneider and the anonomous referee for their invaluable comments on this manuscript. The W.M. Keck Observatory is operated as a scientific partnership between the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration.  It was made possible by the generous financial support of the W. M. Keck Foundation. LML graciously acknowledges support from a Carnegie Fellowship. Support for this work was also provided, in part, by NASA through grant number GO-06000.01-94A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.  \\clearpage"
    },
    "9804/astro-ph9804004_arXiv.txt": {
        "abstract": "s{We review estimates of small scale fluctuations due to extragalactic point sources in the Planck Surveyor frequency bands. While our undestanding of the spectral and evolutionary properties of these sources is far from complete, conservative estimates allow us to confidently conclude that, in the frequency range 100--200 GHz, their contaminating effect is well below the expected anisotropy level of the cosmic microwave background (CMB), down to angular scales of at least $\\simeq 10'$. Hence, an accurate subtraction of foreground fluctuations is not critical for the determination of the CMB power spectrum up to multipoles $\\ell \\simeq 1000$. In any case, Planck's wide frequency coverage will allow to carefully control foreground contributions. On the other hand, the all sky surveys at 9 frequencies, spanning the range 30--900 GHz, will be unique  in providing complete samples comprising from several hundreds to many thousands of extragalactic sources, selected in an essentially unexplored frequency region. New classes of sources may be revealed in these data. The familiar ``flat''-spectrum radio sources should show spectral features carrying essential information on their physical properties. Crucial information will be provided to understand the nature of radio sources with strongly inverted spectra. Scenarios for the cosmological evolution of galaxies will be extensively tested.} ",
        "introduction": "The multifrequency all-sky maps produced by the Planck Surveyor mission will comprise, in addition to anisotropies which are outgrowths of primordial fluctuations, and whose precision measurements are the main goal of the mission, astrophysical foregrounds, the most important of which, over the frequency range of interest, are those due to emissions in our own Galaxy and to extragalactic radio and mm/sub-mm sources.  We deal here with extragalactic sources, which may be a major limiting factor for experiments, like Planck, aimed at accurately determining the cosmic microwave Background (CMB) power spectrum $C_\\ell$ up to multipoles $\\ell \\sim 2000$, corresponding to angular scales $\\theta \\sim 5'$. In fact, a Poisson distribution of sources produce a white noise power spectrum with the same power in all multipoles~\\cite{Tegmark}, so that their contribution to fluctuations in a unit logarithmic multipole interval increases with $\\ell$ as $\\ell(\\ell +1)C_\\ell \\propto \\ell^2$ (for large values of $\\ell$), while, at least for the standard inflationary models, which are consistent with the available anisotropy detections, the function $\\ell(\\ell +1)C_\\ell$ yielded by primordial CMB fluctuations is approximately constant for $\\ell \\lsim 100$, then oscillates and finally decreases quasi exponentially for $\\ell \\gsim 1000$ ($\\theta \\lsim 10'$). Hence confusion noise due to discrete sources will dominate at small enough angular scales.  In \\S$\\,$2 we summarize the limitations set by fluctuations due to extragalactic sources on Planck measurements of primordial CMB anisotropies. On the other hand, the multifrequency all sky surveys carried out by the Planck Surveyor mission will provide a very rich database for astrophysical studies; their impact on investigations of physical and evolutionary properties of different classes of extragalactic sources is briefly outlined in \\S$\\,$3. Our main conclusions are presented in \\S$\\,$4. ",
        "conclusions": "Luckily enough,  both for galaxies and active galactic nuclei, the crossover between the radio and the dust emission components, determining a minimum in the spectral energy distribution, is roughly coincident with the CMB intensity peak. The dust temperature tends to be higher for bright distant objects, moving the minimum to higher frequencies in the rest frame and thus partially compensating for the effect of redshift. This situation makes the mm region ideal for mapping primordial anisotropies.  Although our understanding of foregrounds at Planck frequencies is far from complete, estimates using worst-case parameters in extrapolating existing measurements to Planck frequencies or angular scales, allow us to safely conclude that, in the frequency range 100-200 GHz, the foreground fluctuations, which are dominated, on small scales ($\\theta \\lsim 30'$), by extragalactic sources, are well below the expected amplitude of CMB anisotropies over much of the sky. Hence, the removal of foreground contamination is not critical for accurate determinations of the power spectrum of CMB anisotropies up to multipoles of at least $\\ell \\sim 1000$.  On the other hand, while only a small fraction of high Galactic latitude pixels are strongly contaminated by astrophysical foregrounds, the Planck surveys at 9 frequencies will provide sufficiently rich complete samples for astrophysical studies. Spectral information will be provided for ``flat''-spectrum radio sources (compact radio galaxies, radio loud QSOs, BL Lacs, blazars) over a frequency region where spectral features carrying essential information on their physical conditions show up (breaks due to energy losses of relativistic electrons, self-absorption turnovers of flaring components, ...). Planck surveys will be unique in providing complete samples of bright radio sources with inverted spectra, essentially undetectable in radio-frequency surveys. Important classes of sources of this kind are GHz peaked spectrum sources, which may be the youngest stages of radio source evolution and may thus provide insight into the genesis of radio sources, and advection dominated sources, corresponding to final stages of accretion in giant elliptical galaxies hosting a massive black hole. The high frequency Planck channels will detect thousands of dusty galaxies, a large fraction of which at substantial redshifts, allowing to extensively test scenarios for galaxy evolution. The increasing evidence that a large, and perhaps dominant fraction, of star formation at high redshifts may be hidden by dust, makes far-IR to sub-mm surveys an essential complement to optical data."
    },
    "9804/astro-ph9804142_arXiv.txt": {
        "abstract": "Using the IRAM 30-m telescope, we observed the supernova remnant 3C~391 (G31.9+0.0) and its surroundings in the \\cotwo, \\hcop, \\cstwo, \\csthree, and \\csfive\\ lines. The ambient molecular gas at the distance (9 kpc) of the remnant comprises a giant molecular cloud whose edge is closely parallel to a ridge of bright non-thermal radio continuum, which evidently delineates the blast-wave into the cloud. We found that in a small (0.6 pc) portion of the radio shell, the molecular line profiles consist of a narrow (2 \\kms) component, plus a very wide ($> 20$ \\kms) component. Both spectral components peak within $20^{\\prime\\prime}$ of a previously-detected OH 1720 MHz maser. We name this source 3C~391:BML (broad molecular line); it provides a new laboratory, similar to IC 443 but on a larger scale, to study shock interactions with dense molecular gas. The wide spectral component is relatively brighter in the higher-excitation lines. We interpret the wide spectral component as post-shock gas, either smoothly accelerated or partially dissociated and reformed behind the shock. The narrow component is either the pre-shock gas or cold gas reformed behind a fully dissociative shock. Using the 3 observed CS lines, we measured the temperature, CS column density, and H$_2$ volume density in a dense clump in the parent molecular cloud as well as the wide-line and narrow-line portions of the shocked clump. The physical conditions of the narrow-line gas are comparable to the highest-density clumps in the giant molecular cloud, while the wide-line gas is {\\it both} warmer and denser. The mass of compressed gas in 3C~391:BML is high enough that its self-gravity is significant, and eventually it could form one or several stars. ",
        "introduction": "Supernovae are thought to be the source of kinetic energy of the interstellar medium, keeping the gas in motion and returning material from dense molecular clouds into the more diffuse interstellar medium and the galactic halo.  When a massive star ends its life in a supernova explosion, it often does so in the vicinity of the molecular cloud in which it was born, as is evidenced by the close correspondence of OB associations and giant H~II regions in spiral arms (\\cite{elmlad77}).  Despite the expected close association between Type II supernovae and molecular clouds, very few cases of supernova-molecular cloud (SN-MC) interaction are known or suspected. The blast wave from a supernova within or near the edge of a cloud will progress rapidly through the inter-clump medium and drive slower shocks into dense clumps. Multiple reflections of high-energy charged particles within the complicated magnetic field of an SN-MC interaction are a possible source of cosmic rays, which will permeate the entire region (\\cite{chevalier77}; \\cite{esposito96}). The thermal radiation from the remnant interior (mainly X-rays), cosmic rays and their secondary gamma rays, and direct impact of the blast wave onto clumps should visibly perturb the excitation, chemistry, and dynamics of the parent molecular cloud for at least the $\\sim 10^5$ year period during which the SN blast wave is most powerful.  So far, the only well-known case of an SN-MC interaction is IC 443, where molecular lines have been detected with FWHM $\\sim$ 20 km s$^{-1}$ (much wider than the lines from nearby, un-shocked gas), and from energy levels far above the ground state (\\cite{white87}; van Dishoeck, Jansen, \\& Phillips 1993; \\cite{wang92}). Other remnants, including W~28, CTB 109, Kes 79, and W~51C have been suggested as SN-MC interactions based on their proximity to molecular clouds, wide molecular lines, or both (\\cite{woot77}; \\cite{woot81}; \\cite{tatematsu90}; \\cite{green92}; \\cite{koo97a}; \\cite{koo97b}). In the case of W~28, W~44 and 3C~391, 1720 MHz OH emission has been detected from many small spots, with brightness temperatures so high that they must be masers; these masers are thought to be collisionally excited and they strongly suggest the presence of SN-MC interactions (Frail, Goss, \\& Slysh 1994; \\cite{frail96}). 3C~391 is one of the brightest radio supernova remnants, and high-resolution radio images suggest a `break-out' morphology due to an explosion near the edge of a molecular cloud (\\cite{rm93}). The X-ray emission from 3C~391 peaks in its interior and has a thermal spectrum (\\cite{rp96}), characteristic of a newly-defined class of supernova remnants, called `mixed-morphology' remnants, whose nature has been linked to interaction with a strongly inhomogeneous pre-shock interstellar medium (\\cite{rp98}). A recent map of the \\coone\\ emission in the vicinity of 3C~391 revealed a giant molecular cloud that is precisely parallel to the bright ridge of radio emission, confirming that its `break-out' radio morphology is indeed due to the strong density contrast between the molecular cloud to the northwest and the relatively empty regions elsewhere (Wilner, Reynolds, \\& Moffett 1998). The work described in this paper is part of our recently-initiated campaign to search for and characterize SN-MC interactions in the mixed-morphology supernova remnants. Our first result was the detection of bright [O~I] 63 $\\mu$m and dust emission from 3C~391, showing that the blast-wave into the molecular gas is radiative, and the SN-MC interaction is a significant energy loss for the remnant, although it remains globally adiabatic (\\cite{reach96}). In this paper, we present new observations of molecular emissions from 3C~391, designed to search for the effects of the SN-MC interaction on the molecular cloud, using millimeter-wave observations at high angular resolution and several transitions requiring a range of physical conditions for excitation. Throughout this paper, we assume a distance to 3C~391 of 9~kpc, which is based on the comparison of H~I 21-cm emission and absorption line profiles (\\cite{radakrish}) and the H$_2$CO absorption line at 96 \\kms\\ (\\cite{downes}); our adopted distance is consistent with that adopted by others for this remnant (\\cite{rm93}). ",
        "conclusions": "We observed the entire supernova remnant 3C~391 in millimeter lines of CS, CO, and HCO$^+$. The lower-excitation lines reveal a giant molecular cloud to the northwest of the remnant, explaining the `break-out' morphology of the radio emission. The interactions between the blast wave and a very dense molecular clump was found within a small ($50^{\\prime\\prime}$) region that we call 3C~391:BML. A wide component (FWHM 25 \\kms) and a narrow component (FWHM 2 \\kms) both peak at 3C~391:BML, which is coincident with an OH 1720 MHz maser. The excitation of the wide molecular lines requires both high gas temperature ($> 100$ K) and density ($\\sim 3\\times 10^5$ cm$^{-3}$). The narrow-line region require somewhat lower density and are consistent with much lower ($\\sim 20$ K) temperatures. We identified a clump in the parent molecular cloud with properties similar to the narrow-line region. Therefore, the 3C~391:BML clump was similar to the highest-density clumps in the parent molecular cloud, and it is currently being shocked. The brightness of the wide \\cotwo\\ line from 3C~391:BML is consistent with C-type molecular shocks with $10^4 < n_0 <  10^5$ cm$^{-3}$ and $10 < v_s < 50$ \\kms, or J-type shocks with $n_0\\sim 10^3$ and $v_s \\sim 100$ \\kms. The pressure in the shocked clump is much higher than the estimated ram pressure of the remnant, possibly because of its self-gravity. If so, this clump is a likely site of triggered star formation. A widespread interaction of 3C~391 with molecular gas is evidenced by the distribution of CO and CS lines with central velocities offset from that of the parent cloud by 10 to 15 \\kms; these velocities correspond to those of the two OH masers. The interaction comprises nearly an entire hemisphere of the remnant, making 3C~391 a `CO shell' remnant.  {\\bf Acknowledgment} We thank Hans Ungerechts at IRAM for helping us get started on the IRAM telescope and David Wilner and Steve Reynolds for sharing early results of their observations. WTR thanks the Commissariat d'Energie Atomique, in Saclay, France for hospitality and computing power during part of the data analysis. We thank Bon-Chul Koo for his comments and support. The research described in this paper was carried out in part by the California Institute of Technology, under a contract with the National Aeronautics and Space Administration.  \\clearpage  \\begin{deluxetable}{lllllll} \\footnotesize \\tablecaption{Observed spectral lines and Telescope parameters\\label{tab:telparams}} \\tablewidth{0pt} \\tablehead{ \\colhead{transition} & \\colhead{frequency} & \\colhead{$T_{sys}$\\tablenotemark{a}} & \\colhead{$\\eta_{mb}$} &  \\colhead{beam} & \\colhead{$\\delta v$\\tablenotemark{b}}  &    \\colhead{$\\Delta v$\\tablenotemark{c}} \\\\ & \\colhead{(GHz)}     & \\colhead{(K)} &                           & \\colhead{($^{\\prime\\prime}$)} & \\colhead{(\\kms)}  & \\colhead{(\\kms)}  } \\startdata HCO$^+$($1\\rightarrow 0$) &  89.1885 & 310 & 0.82 & 27 & 0.22 & 430\\nl CS($2\\rightarrow 1$)      &  98.9798 & 250 & 0.76 & 24 & 0.24 & 430\\nl CS($3\\rightarrow 2$)      & 146.9690 & 240 & 0.58 & 16 & 0.65 & 290\\nl CS($5\\rightarrow 4$)      & 244.9356 & 490 & 0.43 & 10 & 1.2  & 310\\nl CO($2\\rightarrow 1$)      & 230.5380 & 720 & 0.45 & 10 & 1.3  & 330\\nl \\enddata \\tablenotetext{a}{typical system temperature for observations presented in this paper} \\tablenotetext{b}{velocity resolution} \\tablenotetext{c}{velocity coverage} \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{llll} \\tablecolumns{4} \\footnotesize \\tablecaption{Measured properties of spectral lines\\tablenotemark{a}\\label{tab:spectab}} \\tablewidth{0pt} \\tablehead{ \\colhead{transition} & \\colhead{$T_{mb}$} & \\colhead{$V_{LSR}$} &  \\colhead{$\\Delta V$ (FWHM)} \\\\ & \\colhead{(K)} & \\colhead{(\\kms)}  & \\colhead{(\\kms)}  } \\startdata  \\cutinhead{wide-line position in shocked clump $(-40^{\\prime\\prime},-50^{\\prime\\prime})$\\tablenotemark{b}} HCO$^+$($1\\rightarrow 0$) & 0.91 & 111.4 & 25.5 \\nl CS($2\\rightarrow 1$)      & 0.32 & 108.9 & 19.0 \\nl & 0.24 & 104.2 &  1.4 \\nl CS($3\\rightarrow 2$)\t  & 0.37 & 108.9 & 20.0 \\nl CS($5\\rightarrow 4$)\t  & 0.35 & 108.5 & 15.6 \\nl CO($2\\rightarrow 1$)\t  & 17.6 & 111.1 & 22.5 \\nl & 15.1 & 104.2 &  2.6 \\nl  \\cutinhead{narrow-line position in shocked clump $(-10^{\\prime\\prime},-85^{\\prime\\prime})$\\tablenotemark{b} } HCO$^+$($1\\rightarrow 0$) &  1.0 & 105.3 &  1.7 \\nl & 0.38 & 102.6 & 14.4 \\nl CS($2\\rightarrow 1$)      & 0.90 & 105.5 &  1.0 \\nl & 0.23: & 103.7 &  6.0 \\nl CS($3\\rightarrow 2$)\t  & 0.55 & 105.4 &  1.2 \\nl & 0.23 & 103.7 &  6.0 \\nl CS($5\\rightarrow 4$)\t  & $<0.13$ & & \\nl CO($2\\rightarrow 1$)\t  & 11.6 & 105.5 & 1.2 \\nl &  9.7 & 104.8 & 3.7 \\nl  \\cutinhead{radio ridge $(-130^{\\prime\\prime},+90^{\\prime\\prime})$\\tablenotemark{c} } HCO$^+$($1\\rightarrow 0$) & 0.22: &  96.4 &  2.4 \\nl CS($2\\rightarrow 1$)      &  $<0.03$ & & \\nl CS($3\\rightarrow 2$)      &  $<0.04$ & & \\nl CS($5\\rightarrow 4$)      &  $<0.13$ & & \\nl CO($2\\rightarrow 1$)\t  &  8.9 &  96.5 &  4.0 \\nl &  3.3 & 102.0 &  2.1 \\nl &  3.6 & 107.0 &  3.8 \\nl  \\cutinhead{clump in parent cloud $(-130^{\\prime\\prime},+240^{\\prime\\prime})$\\tablenotemark{c} } HCO$^+$($1\\rightarrow 0$) &  0.32 &  96.1 & 7.3: \\nl CS($2\\rightarrow 1$)      &  0.55 &  96.7 & 3.8 \\nl CS($3\\rightarrow 2$)\t  &  0.27 &  96.5 & 3.6 \\nl CS($5\\rightarrow 4$)\t  & $<0.13$ & & \\nl CO($2\\rightarrow 1$)\t  & 12.0 &  96.0 & 5.4 \\nl &  8.4 & 107.9 & 4.0 \\nl \\enddata \\tablenotetext{a}{only components between 90 and 140 \\kms\\ are listed} \\tablenotetext{b}{spectra averaged within $11^{\\prime\\prime}$ radius} \\tablenotetext{c}{spectra averaged within $21^{\\prime\\prime}$ radius} \\end{deluxetable}  \\clearpage"
    },
    "9804/astro-ph9804156_arXiv.txt": {
        "abstract": "The scattering diameters of \\sgra\\ and several nearby OH masers ($\\approx 1\\arcsec$ at 1~GHz) indicate that a region of enhanced scattering is along the line of sight to the Galactic center.  The scattering diameter of an extragalactic source seen through this scattering region will be larger by the ratio of the Sun-GC distance to the GC-scattering region separation.  This ratio could be a factor of a few, if the scattering region is far from the GC and only a random superposition with it, to more than 100, if the scattering region is within the \\hbox{GC}.  We have used the VLA to survey 10 (11) fields at 20~cm (6~cm) that are between 7\\arcmin\\ and 137\\arcmin\\ from \\sgra.  Our objective was to identify extragalactic sources and measure their scattering diameters so as to constrain the GC-scattering region separation.  In order to find sources within these fields, we have employed pdf\\clean, a source detection algorithm in which sources are identified in an image by comparing the intensity histogram of the image to that expected from a noise-only image.  We found over 100 sources, with the faintest sources being approximately 3~mJy.  The average number of sources per field is approximately 10, though fields close to \\sgra\\ tend to contain fewer sources.  In a companion paper we combine our survey with previous observations of the GC, and we assess the likelihood that the scattering region is so close to the GC that the resulting scattering diameters cause extragalactic sources to be resolved out by our observations.  A number of Galactic sources is included in our source catalog.  We discuss the double-lobed source 1LC~359.872$+$0.178, potentially an X-ray quiet version of 1E~1740.7$-$2942, a shell-like structure with a central point source, and a possible radio transient. ",
        "introduction": "\\label{sec:gc.intro}  If viewed through a plasma containing density fluctuations, an otherwise unresolved source will have a visibility, as measured by an interferometer of baseline length~$b$, of \\begin{equation} V(b) = \\exp\\left[-\\frac{1}{2}D_\\phi(b)\\right]. \\label{eqn:visibility} \\end{equation} The phase structure function, $D_\\phi(b) \\equiv \\langle[\\phi(0)-\\phi(b)]^2\\rangle$, is a measure of the phase perturbations, on a length scale~$b$, imposed on a propagating electromagnetic wave by fluctuations in the electron density.  For a plane wave impinging on this scattering region \\begin{equation} D_\\phi(b) = 8\\pi^2r_{\\mathrm{e}}^2\\lambda^2\\int_0^D dz\\,\\int_0^\\infty dq\\,q[1 - J_0(bq)]P_{\\delne}(q, z), \\label{eqn:structurefunction1} \\end{equation} where $r_{\\mathrm{e}}$ is the classical electron radius, $J_0(x)$ is the zeroth-order Bessel function, $P_{\\delne}$ is the spatial spectrum of the density fluctuations, and the integral over $z$ is taken \\emph{from the source to the observer}.  If the source of radiation is close to or embedded within the scattering medium, so that the medium is illuminated by spherical wavefronts, the argument of the Bessel function is $bq(z/D)$ (\\cite{i78}); the factor~$z/D$ accounts for the divergence of spherical waves.  The apparent angular diameter of the source is determined by the width of the visibility function, and, hence, by how quickly $D_\\phi(b)$ decreases as a function of~$b$.  Since $z/D < 1$, the difference in the form of the phase structure function for plane and spherical wavefronts means that sources close to the medium will show smaller angular diameters than those far from it.  Hence, by comparing the scattering diameters of Galactic and extragalactic sources along similar lines of sight, one can constrain the \\emph{radial} location of the scattering material.  Toward the Galactic center, the observed diameter of \\sgra\\ scales as $\\lambda^2$ over the wavelength range 30~cm to~3~mm (Davies, Walsh, \\& Booth~1976; \\cite{rogersetal94}), as expected if very strong interstellar scattering from microstructure in the electron density determines the observed diameter.  Maser spots in OH/IR stars within 25\\arcmin\\ of \\sgra\\ also show enhanced angular broadening (\\cite{vfcd92}; \\cite{fdcv94}).  The scattering disks of \\sgra\\ and many of the OH masers are observed to be anisotropic as well (\\cite{vfcd92}; \\cite{bzkrml93}; \\cite{krichbaumetal93}; \\cite{fdcv94}; \\cite{y-zcwmr94}); in the case of \\sgra, its scattering disk is anisotropic at least over the wavelength range 21~cm to 7~mm. These observations indicate that a region of enhanced scattering with an angular extent of at least 25\\arcmin\\ (60~pc at 8.5~kpc) is along the line of sight to \\sgra.  At 1~GHz the level of angular broadening produced by this scattering region is roughly 10 times greater than that predicted by a recent model for the distribution of free electrons in the Galaxy (Taylor \\& Cordes~1993, hereinafter \\cite{tc93}), even though this model includes a general enhancement of scattering toward the inner Galaxy.  Because all of the sources observed through this region have thus far been Galactic sources, with (presumably) approximately the same location (i.e., in the Galactic center), the radial location of the scattering region is unconstrained.  The scattering region could be local to the Galactic center, within approximately 100~pc from the Galactic center---which we refer to as the GC model---or the region could be a random superposition and more than 1~kpc from the GC---which we refer to as the RS model.  In the GC model, the region would be a site of excess scattering, and presumably arises from processes unique to the GC; in the RS model, the level of scattering in the region would be high, but not unusually so.  Previous estimates for the location of the scattering region have ranged from 10~pc to 3~kpc.  Ozernoi \\& Shisov~(1977) concluded that an ``unrealistic'' level of turbulence is implied unless the region is within 10~pc of the \\hbox{GC}.  The level of turbulence they considered unrealistic, however, namely $\\sqrt{\\langle n_{\\mathrm{e}}^2\\rangle}/\\langle n_{\\mathrm{e}}\\rangle \\sim 1$, does appear to occur elsewhere in the interstellar medium (\\cite{s91}). Further, van~Langevelde et al.~(1992) showed that the free-free absorption toward \\sgra\\ would be excessive unless the scattering region was at least 0.85~kpc from the GC, though suitable adjustment of free parameters (outer scale and electron temperature) can decrease the limit to 0.03~kpc.  With the free-free absorption they also placed an upper limit on the region's distance from the GC of 3~kpc. Although the GC model is attractive for phenomenological reasons, other sites of enhanced interstellar scattering are found throughout the Galaxy (e.g., NGC~6634, \\cite{mrgb90}; Cyg~X-3, \\cite{mmrj95}) and the mean free path for encountering such a region is approximately 8~kpc (\\cite{cwfsr91}).  Identifying the location of the scattering region may provide clues to the origin of the scattering.  The density fluctuations responsible for interstellar scattering are believed to be generated by velocity or magnetic field fluctuations (\\cite{h84}, 1986; Montgomery, Brown, \\& Matthaeus~1987; \\cite{s91}; \\cite{sg94}; \\cite{gs95}).  Velocity or magnetic field fluctuations are also a natural means for inducing anisotropy in the density fluctuations and thereby in the scattering disks.  If this supposition is correct, the amplitude of the density fluctuations may provide a measure of the coupling between the density and velocity or magnetic field fluctuations or, more generally, provide information about the small-scale velocity or magnetic field in the scattering region.  However, because the radial location of the scattering region is unconstrained, relevant quantities, e.g., the rms density, are uncertain by a factor of $\\delgc/\\dgc$, where $\\delgc$ is the GC-scattering region separation and $\\dgc$ is the GC-Sun distance.  Observations of extragalactic sources viewed through the scattering region could constrain $\\delgc$; however, few extragalactic sources have been identified toward the \\hbox{GC}.  The two sources closest to \\sgra\\ are B1739$-$298 (\\cite{dkvgh83}) and GPSR~0.539$+$0.263 (Bartel~1994, private communication), which are 48\\arcmin\\ and 40\\arcmin\\ from \\sgra, respectively.  Neither of these is within the region of enhanced scattering defined by the OH masers.  This paper reports VLA and VLBA observations of potential extragalactic sources seen through the \\hbox{GC}. Section~\\ref{sec:gc.observe} describes the observations and data reduction, Section~\\ref{sec:catalog} discusses the identification of potential extragalactic sources and presents the catalog of sources, Section~\\ref{sec:sources} discusses certain Galactic sources found in our VLA survey, and Section~\\ref{sec:gc.conclude} discusses our results and presents our conclusions.  A companion paper (Lazio \\& Cordes~1998, hereinafter \\cite{lc98}) combines the results of this paper with the previous observations of OH and \\hoh\\ masers and free-free emission in a likelihood analysis that constrains the angular extent and radial location of the scattering region. \\cite{lc98} also discusses the physical conditions inside the scattering region. ",
        "conclusions": "\\label{sec:gc.conclude}  This paper has reported the results of a program to identify and obtain scattering diameters for extragalactic sources seen through the Galactic center scattering region.  Because they are located far behind the GC, the scattering diameters of extragalactic sources, when compared to the scattering diameter of GC sources such as \\sgra, can constrain the \\emph{radial} location of the scattering region, viz.\\ equation~(\\ref{eqn:xgalsize}) and Figure~\\ref{fig:xgalsize}.  Using the VLA we observed 10 (11) fields at~20~cm (6~cm) containing 15 suspected extragalactic sources.  We increased our catalog of sources to well over 100 through the use of pdf\\clean: The intensity histogram of the primary beam was used to identify positive brightness image pixels that produced deviations from the shape of the expected noise-only histogram.  We found approximately 10 sources per field.  Follow-up VLBI observations on a subset of these sources have determined the scattering diameters for two heavily scattered extragalactic sources.  Their diameters are too small, by factors of 4--10, for them to be seen through the scattering region in front of \\sgra.  However, they can be used, in conjunction with the heavily scattered masers, to set constraints on the angular extent of the region.  Our fields show a paucity of sources near \\sgra; a previous survey with more uniform sky coverage, but at a lower sensitivity also shows a paucity.  Such a deficit could arise if the scattering toward the GC is so severe that our (and previous) observations resolve out extragalactic sources.  The sources reported here are combined with angular broadening measurements of \\sgra\\ and OH masers and free-free emission and absorption measurements from the literature.  These data are then used in a likelihood analysis to determine the model parameters of the GC scattering region (\\cite{lc98})."
    },
    "9804/astro-ph9804010_arXiv.txt": {
        "abstract": "In this paper, the third of a series dedicated to the investigation of the nuclear properties of spiral galaxies, we have {\\it (i)} modelled the {\\tt WFPC2} F606W nuclear surface brightness profiles of 41 spiral galaxies presented in Carollo et al.\\ 1997c, 1998 with the analytical law introduced by Lauer et al.\\ 1995, and {\\it (ii)} deprojected these surface brightness profiles and their analytical fits, so as to estimate the nuclear stellar densities of bulges of spiral galaxies. We find that the nuclear stellar cusps (quantified by the average logarithmic slope of the surface brightness profiles within 0.1$''$-0.5$''$) are significantly different for $R^{1/4}$-law and exponential bulges. The former have nuclear properties similar to those of early-type galaxies, i.e.  similar values of nuclear cusps for comparable luminosities, and increasingly steeper stellar cusps with decreasing luminosity. By contrast, exponential bulges have (underlying the light contribution from photometrically distinct, central compact sources) comparative shallower stellar cusps, and likely lower nuclear densities, than $R^{1/4}$-law bulges. ",
        "introduction": "The galactic nuclei are the repositories of low angular momentum material sunk to the centers over the lifetime of the parent systems. Therefore, they are likely to hold answers to important questions related with the origin of structure in the parent galaxies. In this perspective, establishing the demographics of galactic nuclei along the entire Hubble sequence lies at the heart of our understanding of the complex process of galaxy formation and evolution.  Observations of nearby ellipticals and lenticulars with the Faint Object Camera ({\\tt FOC}), the Wide Field Planetary Camera ({\\tt WF/PC}) and the Wide Field Planetary Camera-2 ({\\tt WFPC2}) aboard the {\\it Hubble Space Telescope} (HST) have revealed that the nuclei of these galaxies are complex environments (e.g., Crane et al.\\ 1993; Jaffe et al.\\ 1994; Lauer et al.\\ 1995, hereafter L95; Forbes et al.\\ 1995; Carollo et al.\\ 1997a, hereafter C97a; Carollo et al.\\ 1997b, hereafter C97b; Faber et al.\\ 1997, hereafter F97).  They show surface brightness profiles that increase down to the innermost point measurable at HST resolution, i.e., $I(r) \\propto r^{-\\gamma}$ as $r \\rightarrow 0$ (where $I(r)$ is the surface brightness at the radius $r$, and $\\gamma > 0$); furthermore, several galaxies host stellar and gaseous disks, unresolved nuclear spikes, double nuclei. These inner features might possibly be related to the presence of massive black holes (e.g., Lauer et al.\\ 1996).  By contrast, much is still to be learned about the nuclear properties of nearby spiral galaxies at HST resolution. F97 found that the surface brightness profiles of the three Sa-Sb bulges present in their sample show a behaviour similar to that of early-type spheroidals of comparable luminosity. The same result was found by Phillips et al.\\ (1996) for the three spirals of type earlier than Sc contained in their {\\tt WF/PC} F555W sample of 20 disk galaxies. Furthermore, Phillips et al.\\ found that later type spirals show instead (almost) flat nuclear profiles, and suggested that the nuclear properties of disk galaxies are more closely related to those of nucleated dwarf galaxies than to those of elliptical galaxies. Further exploration is necessary to assess how the nuclear properties scale with the properties of the spheroidal component. This is likely to provide feedback information about the epoch and processes of nucleus, bulge, and, ultimately, galaxy formation.  In order to address these issues, we have performed a {\\tt WFPC2} snapshot survey in the F606W filter of the nuclei of 107 (mostly Sa to Sc) disk galaxies. In paper I (Carollo et al.\\ 1997c) and paper II (Carollo et al.\\ 1998) we have presented the 75 targets imaged so far within our program. Our analysis shows that bulge-like structures are present in most of the galaxies.  While in some cases these are ``classical'', smooth, featureless $R^{1/4}$-law bulges, in others they are better fitted by an exponential profile (see also Courteau, de Jong \\& Broeils 1996, and references therein, for similar results). The exponential bulges include two classes of objects: {\\it (i)} dwarf-looking systems, whose surface brightness profiles within $\\approx 15''$ are well fitted by a single exponential. These galaxies are strongly bulge-dominated; their surrounding, faint regions (``disk/halo'') show no signs of spiral arms, and  have typically a quiescent, i.e. non star forming, appearance.  {\\it (ii)} Small exponential bulges embedded in dominant (spiral-armed/star-forming) surrounding disks, i.e., the inner exponential structures of double-exponential fits to the surface brightness profiles within $\\approx 15''$.  The exponential bulges as-a-class are statistically fainter than the featurless, smooth $R^{1/4}$ bulges, for constant disk luminosity and Hubble type. Resolved, central compact sources are found in most of the exponential bulges; the hosts of central compact sources often contain a barred structure. The nature of these compact source, and in particular their relation with e.g., star clusters and Seyfert 2 nuclei, is discussed in Carollo (1998).  In this paper, the third of the series, we investigate the relation between the nuclear structure of spiral galaxies and the physical properties of inner disks and/or bulges.  In particular, we {\\it (i)} present the results of modeling the nuclear surface brightness profiles with the analytical law introduced by L95 (for the 43 galaxies of paper I and II for which we could perform the measurements), {\\it (ii)} deconvolve the surface brightness profiles and their analytical fits in order to estimate the nuclear stellar densities, {\\it (iii)} study the nuclear properties as a function of the global properties discussed in papers I and II (e.g., $R^{1/4}$ against exponential bulges), and {\\it (iv)} compare the nuclear properties of our sample with those observed in early-type galaxies. The paper is organized as follows. In section 2 we briefly summarize the properties of the sample, the data used in our investigation, the procedure adopted for the data reduction, and the steps performed to derive the surface brightness profiles. In section 3 we present the results of the analytical fits applied to the surface brightness profiles, and of deconvolving data and models in spherical symmetry. In section 4 we investigate the dependence of the nuclear properties on global galactic properties. We conclude in section 5. ",
        "conclusions": "In this paper we have investigated the relation between the nuclear structure of spiral galaxies and the physical properties of their bulges. In particular, we have {\\it (i)} modelled the {\\tt WFPC2} F606W nuclear surface brightness profiles of 41 spiral galaxies with the analytical law introduced by L95 (data from papers I and II), and {\\it (ii)} deconvolved the surface brightness profiles and their analytical fits in order to estimate the nuclear stellar densities of disk galaxies.  Our main result is that $R^{1/4}$-law bulges and exponential bulges have significantly different nuclear stellar cusps and densities. Specifically, $R^{1/4}$-law bulges have steep stellar cusps which steepen with decreasing luminosity; furthermore, their stellar cusp slopes and densities are similar in values to those of early-type systems of comparable luminosity. By contrast, in exponential bulges, the inward extrapolations underlying the light from the compact sources which sit in their very centers imply rather shallow stellar cusps and, very likely, relatively low nuclear stellar densities.    \\bigskip \\bigskip  \\noindent {\\bf Acknowledgements}  We heartly thank Tim Heckman and Colin Norman for helpful discussions, and the anonymous referee for constructive comments to an earlier version of this paper. CMC is supported by NASA through the grant HF-1079.01-96a awarded by the Space Telescope Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under contract NAS 5-26555. This research has been partially funded by grant GO-06359.01-95A awarded by STScI, and has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, Caltech, under contract with NASA.   \\bigskip \\bigskip  \\noindent \\begin{center}{{\\bf Appendix A.} Details on the Analytical Fits}\\end{center}  \\medskip  We used the same fitting procedure described in C97a; we addess to this reference for further information. This was carried out in two steps: (1) The first step isolated primary from secondary minima of $\\chi^2$. The $\\chi^2$ values were computed on a grid of points uniformly distributed on a wide hypercube in parameter space (with dimension equal to the number of free parameters). Once a minimum value was found, a new, smaller, hypercube was placed on that location, and the procedure iterated. (2) The minimum value found on the hypercube was then used as starting point to initialize a downhill simplex minimization. We tested the procedure on simulated data, and verified that it recovered the initial values with high accuracy.  We accepted as final the fits associated with the absolute minimum of $\\chi^2$.  Deconvolutions of {\\tt WF/PC} data have been proven to be very reliable (e.g., L95). However, in our analysis, we chose not to apply any deconvolution to the post-refurbishment {\\tt WFPC2} images, and to correct for PSF-blurring while modeling the light profiles. Therefore, the models were convolved with the appropriate PSFs before being compared to the data.  Since pointlike sources with adequate S/N located near to the nuclei were not available for most of the galaxies, we computed the PSFs by running Tinytim (Krist 1992).  Focus drifts and breathing modify the PSF profile and affect the flux within a 1 PC pixel radius up to 10\\% (and within 5 PC pixels up to 5\\%; Suchkov \\& Casertano 1997).  Therefore, the simulated PSFs obtained by construction at the nominal focus position are in principle of similar quality than PSFs derived from archival stars.  Furthermore, our approach of convolving the models rather than deconvolving the data minimizes the effects of using a possibly non-perfect PSF.     \\bigskip \\bigskip  \\noindent \\begin{center}{{\\bf Appendix B.} Classifying $R^{1/4}$-law or exponential bulges outside $1''$}\\end{center}  \\medskip  In order to ensure that the classification of a bulge as an $R^{1/4}$-law or an exponential structure is valid on a radial range entirely different from that used in the derivation of $\\langle \\gamma \\rangle$ (equal to 0.1$''$-0.5$''$), we performed as a test the $R^{1/4}$-law and exponential fits after excluding the data inside the innermost $1''$.  This value is a compromise between a radius large enough to exclude the range where $\\langle \\gamma \\rangle$ is computed, and small enough to still allow the detection of the small, disk-embedded exponential bulges (pentagons in the figures).  As an example, the results of the test are illustrated in Figure 8 for the same four galaxies presented in Figure 6. The solid lines represent either a single exponential (left panels) or a double exponential (right panels) fit; the dashed lines represent either a single $R^{1/4}$-law (left panels) or an $R^{1/4}$-law plus exponential (right panels) fit. Two different scales are used for the abscissa for the galaxies in the left and right panels, consistently with the different scales of their bulge components. An offset of two magnitudes has been applied to ESO482G17 for plotting purposes. There are two important points to note: {\\it (i)} the kind of profile which provides the bulge classification given in papers I and II, i.e. exponential or $R^{1/4}$-law, still provides a better fit to the inner galactic regions, even when the innermost $1''$ is excluded from the fits; {\\it (ii)} the alternative profile with respect to the one that provides the classification given in papers I and II generally provides (not only a worse fit but also) physically meaningless best fit parameters (e.g., for ESO482G17, the $R^{1/4}$-law best fit of Figure 8 has an $R_e\\sim385''$).  We conclude that the distinction between $R^{1/4}$-law and exponential bulges holds in a radial range which excludes the one used to derive $\\langle \\gamma \\rangle$, and that the difference in nuclear cusp slopes $\\langle \\gamma \\rangle$ between $R^{1/4}$-law and exponential bulges has a physical origin. We retained the bulge parameters presented in papers I and II for our discussion, since those fits gave an overall better description of the profiles.  \\bigskip \\bigskip"
    },
    "9804/astro-ph9804226_arXiv.txt": {
        "abstract": "An object discovered during an infrared survey of the field near the quasar B2 0149$+$33 has an emission line at 2.25\\,$\\mu$m that we interpret as H$\\alpha$ at a redshift of 2.43. The K-band image shows two compact components 10\\,kpc apart surrounded by more extended emission over $\\sim 20$\\,kpc.  The H$\\alpha$ emission appears to be extended over $\\sim 15$\\,kpc (2$^{\\prime\\prime}$) in a coarsely sampled (0\\farcs8/pixel) image.  The star formation rate may be as high as 250 -- 1000\\,M$_\\odot$\\,yr$^{-1}$, depending on the extinction.  Alternatively, the line may be powered by an active nucleus, although the probability of serendipitously discovering an AGN in the survey volume is only $\\sim 0.02$.  The increasing number of similar objects reported in the literature indicate that they may be an important, unstudied population in the high redshift universe. ",
        "introduction": "Discovering the properties of galaxies at redshifts greater than one requires techniques that can readily distinguish the high redshift objects from those at lower redshift that predominate in any deep image.  Multi-wavelength approaches are especially fruitful and have been dominated to date by optical and radio surveys.  The use of photometric redshifts based on the strong Lyman break redshifted into the optical band has been one of the most successful of these techniques (Steidel et al. 1996, and references therein).  Methods using infrared images are beginning to uncover objects not easily discovered with optical or radio methods that may, nevertheless, constitute a significant fraction of the high redshift population. These include objects distinguished by unusually red colors (Elston, Rieke, \\& Rieke 1988; Soifer et al. 1994; Graham et al. 1994; Hu \\& Ridgway 1994; Cowie et al. 1994; Dey, Spinrad, \\& Dickinson 1995) and objects with emission lines redshifted to infrared wavelengths (Songaila et al. 1994; Thompson, Mannucci, \\& Beckwith 1996, hereafter TMB96; Malkan, Teplitz, \\& McLean 1995; Bechtold et al. 1997). These objects have been interpreted as elliptical galaxies (Graham et al. 1994; Hu \\& Ridgway 1994; Dunlop et al. 1996), young ``protogalaxies'' undergoing bursts of star formation (Eisenhardt \\& Dickinson 1992; Graham \\& Dey 1996; Malkan, Teplitz, \\& McLean 1996; Yee et al. 1996; Bechtold et al. 1997), or active galactic nuclei (Cowie et al. 1994; Dey, Spinrad, \\& Dickinson 1995).  All of these populations could be significant for cosmology, since the first case implies massive galaxy formation at redshifts greater than $\\sim$3, the second indicates a substantial population of young galaxies that can be discovered only in infrared surveys, and the third implies a population of infrared bright active galactic nuclei (AGN) comparable in number density to the populations of AGN discovered by more traditional methods.  Only a few such objects have been studied in enough detail to reveal redshifts, source morphologies, and colors.  In the course of our survey for emission line galaxies (TMB96), an object was discovered near the quasar B2 0149$+$33 that had an emission line at 2.25\\,$\\mu$m, the same wavelength as the quasar's H$\\alpha$ line at a redshift of 2.43; we call this object TMB 0149 - cK39, or simply cK39. A spectrum between 1.5 and 2.4 $\\mu$m, presented and discussed here, confirmed the presence of an emission line nearly coincident in wavelength with that of the quasar. The object is very red, making it unusual among known, distant galaxies.  This paper describes the results of these observations and suggests that such objects may be common but previously unobservable owing to the lack of optical emission. ",
        "conclusions": "The emission line object, TMB 0149 $-$ cK39, appears to be a pair of galaxies undergoing a merger at redshift of 2.4.  If cK39 derives a substantial part of its luminosity from star formation, the formation rate is as high as 1000 M$_\\odot$ yr$^{-1}$, an exceptionally large value that is rarely seen in other starforming galaxies. On the other hand, the system could contain at least one active nucleus, perhaps with some contribution from star formation.  The individual components appear to be 6 and 7 kpc in extent with the centers separated by 10 kpc, consistent with a merger-induced fuelling of the nuclear activity.  An extinction of A$_{\\rm V} \\sim 1.7^{\\rm m}$ is required to produce the observed R-K color of 5.5, requiring the presence of a significant amount of dust in order to suppress the ultraviolet light and redden the galaxies.  Although gravitational lensing by a foreground galaxy or cluster of galaxies could enhance the brightness of the emission, there is little evidence of such a galaxy or cluster along the line of sight.  The growing number of very red galaxies at high redshift indicates that there are new populations uncovered only in infrared surveys (e.g. Graham \\& Dey 1996; Malkan, Teplitz, \\& McLean 1996; and references therein).  In a separate paper (Thompson et al. in preparation), we present statistics that show extremely red objects (R-K$^\\prime > 6$) have a sky density of order 500 deg$^{-2}$ for K$^\\prime$ $\\leq$19.75, so cK39 may, indeed, be part of a larger population. These results underscore the importance of employing a number of different techniques for exploring the epoch of early star formation and demonstrate that a better understanding of these objects is important to an understanding of galaxy formation in the early universe."
    },
    "9804/astro-ph9804295_arXiv.txt": {
        "abstract": "s{Direct and indirect detection rates of relic neutralinos are reviewed in the framework of the Minimal Supersymmetric Standard Model. The theoretical estimates are compared with the most recent experimental limits from low--background detectors and neutrino telescopes. The properties of neutralino under the hypothesis that preliminary experimental results of the DAMA/NaI Collaboration may be indicative of a yearly modulation effect are examined.}   \\normalsize\\baselineskip=15pt  \\vspace{-5pt} ",
        "introduction": "\\vspace{-5pt}  Supersymmetric theories predict a large number of particles in excess to the Standard Model ones. If the R--parity is conserved, the lightest among all the supersymmetric particles (LSP) must be stable. This feature makes the LSP a dark matter candidate, since this particle can be present today as a relic from the early stages of the evolution of the Universe. Different candidates have been proposed in the framework of supersymmetric theories: the neutralino or the sneutrino\\cite{sneutrino} in gravity mediated models, the gravitino\\cite{gravitino} or some messenger fields in gauge mediated theories\\cite{messenger}, the axino\\cite{axino}, stable non--topological solitons (Q--balls)\\cite{Qballs} or others.  In this paper we will focus on the most promising among all the different candidates, the neutralino, which is defined as the lowest mass linear superposition of photino ($\\tilde \\gamma$), zino ($\\tilde Z$) and the two higgsino fields ($\\tilde H_1^{\\circ}$, $\\tilde H_2^{\\circ}$), i.e. $\\chi \\equiv a_1 \\tilde \\gamma + a_2 \\tilde Z + a_3 \\tilde H_1^{\\circ} + a_4 \\tilde H_2^{\\circ}$. The aim of this review is at providing the latest results on the calculation of different kinds of detection rates of relic neutralinos, in the framework of the Minimal Supersymmetric extension of the Standard Model (MSSM), constrained by the most recent experimental data coming from accelerator physics. We do not discuss here the details of the model, for which we refer to Refs.\\cite{pinning,extending} and to the references quoted therein. We only recall the standard assumptions employed here: i) all trilinear parameters are set to zero except those of the third family, which are unified to a common value $A$; ii) all squarks and sleptons soft--mass parameters are taken as degenerate: $m_{\\tilde l_i} = m_{\\tilde q_i} \\equiv m_0$; iii) the gaugino masses are assumed to unify at $M_{GUT}$, and this implies $M_1= (5/3) \\tan^2 \\theta_W M_2$ at the electroweak scale. After these conditions are applied, the free parameters are: $M_2, \\mu, \\tan\\beta, m_A, m_0, A$. The parameters are varied in the following ranges: $10\\;\\mbox{GeV} \\leq M_2 \\leq  500\\;\\mbox{GeV},\\; 10\\;\\mbox{GeV} \\leq |\\mu| \\leq  500\\;\\mbox{GeV},\\; 65\\;\\mbox{GeV} \\leq m_A \\leq  500\\;\\mbox{GeV},\\; 100\\;\\mbox{GeV} \\leq m_0 \\leq  500\\;\\mbox{GeV},\\; -3 \\leq {\\rm A} \\leq +3,\\; 1.01 \\leq \\tan \\beta \\leq 50$. In our analysis the supersymmetric parameter space is constrained by all the experimental limits on Higgs, neutralino, chargino and sfermion searches at accelerators. Moreover, the constraints due to the $b \\rightarrow s + \\gamma$ process\\cite{LEP} are satisfied. In addition to the experimental limits, we require that the neutralino is the lightest supersymmetric particle. Finally, the regions of the parameter space where the neutralino relic abundance exceeds the cosmological bound, i.e. $\\Omega_{\\chi}h^2 > 1$, are also excluded.  \\vspace{-7pt} ",
        "conclusions": "\\vspace{-10pt} In this paper we have reported the most recent calculations of the direct and indirect detection rates of relic neutralinos in the framework of the Minimal Supersymmetric Standard Model. We have shown that the theoretical estimates of the detection rates may be at the level of the present experimental sensitivities of low--background detectors and neutrino telescopes. For many supersymmetric configurations, and for median values of the astrophysical parameters which enter in the calculations of the detection rates, the predicted signals may already exceed the present experimental bounds. This shows that the different experimental efforts to search for relic particles are potentially able to deeply investigate the possibility that neutralino is a component of the dark matter of the Universe. An interesting preliminary analysis of the DAMA/NaI Collaboration has shown that their data are compatible, at 90\\% C.L., with a modulation signal of the direct detection rate. The features of a neutralino able to satisfy the prerequisites of this signal have been analyzed and it has been shown that many configurations are compatible with a dark matter scenario where the neutralino is the major component, both on galactic and cosmological scales. However, we have to remind here that the occurrence of a possible modulation effect will necessarily require further investigations with much higher statistics. This project is currently under way.  \\vspace{-7pt}"
    },
    "9804/astro-ph9804341_arXiv.txt": {
        "abstract": "Primordial Black Holes (PBH) may have formed from the collapse of cosmic string loops. The spectral shape of the PBH mass spectrum can be determined by the scaling argument for string networks. Limits on the spectral amplitude derived from extragalactic $\\gamma$-ray and galactic $\\gamma$-ray and cosmic ray flux observations as well as constraints from the possible formation of stable black holes remnants are reanalyzed. The new constraints are remarkably close to those derived from the normalization of the cosmic string model to the cosmic microwave background anisotropies. ",
        "introduction": "Cosmic strings (CS) are linear topological defects that are believed to originate during phase transitions in the very early Universe \\cite{VSrev,HKrev,RBrev}. Here, we consider the ``standard\" CS model \\cite{Zel,V81}, according to which the network of linear defects quickly reaches a ``scaling\" solution characterized by having the statistical properties of the string distribution independent of time if all lengths are scaled to the Hubble radius ($R_H = c t$, where $c$ is the speed of light). Cosmic string loops (CSL) are continually formed by the intersection and self-intersection of long CS (infinite CS or CSL with radius of curvature larger than $R_H$). After formation, a loop oscillates due its own tension and slowly decays by emitting gravitational radiation. The initial length of a CSL is $l(t) = \\alpha R_H$, where $\\alpha$ is expected to be $\\sim G \\mu/c^2$. The mass of a CSL is $m(t) = l(t) \\mu$, where $\\mu$ is the mass per unit of length of the string  Since CS also lead to cosmic microwave background (CMB) anisotropies, the string model can be normalized by the recent COBE observations giving the constraint \\cite{LP93,ACSSV} \\be \\label{cmbnorm} G\\mu / c^2 \\leq 1.7(\\pm0.7) \\times 10^{-6} \\ee  Our assumption is that a distribution of PBH was formed by the collapse of a fraction $f$ of CSL \\cite{SH,AP}. Hence, from the observational consequences of a present surviving distribution of PBH we can derive updated constraints on the CS scenario \\cite{jru}. These constraints are important because: {\\it i-)} They may indicate new ways to search for direct signatures from CS; {\\it ii-)} They may provide constraints on CS models with symmetry breaking scale $\\mu^{1/2}$ smaller than $10^{16}$ GeV which are not constrained by CMB and large-scale structure data; and {\\it iii-)} They may provide tighter limits than the CMB on CS models with $G \\mu / c^2 \\sim 10^{-6}$.  Because CS do not dominate the energy density of the Universe, the CS network must lose energy. We derive the rate of CSL production ${{dn_l} \\over {dt}}$ from the conservation of string energy in the ``scaling\" scenario, \\[  \\label{consmass} {\\dot \\rho_{\\infty}} - 2 H \\rho_{\\infty} = - {{dn_l} \\over {dt}} \\alpha \\mu t \\, , \\] where $\\rho_{\\infty} = \\nu \\mu c^{-3} t^{-2}$ is the energy density in long strings and $\\nu$ is proportional to the average number of long strings crossing each Hubble volume.  Hawking \\cite{SH} and Polnarev and Zembowicz \\cite{AP} first postulated that a fraction $f$ of the CSL could collapse within its Schwarzchild radius and then form a BH. More recently, Caldwell and Casper \\cite{CC} have performed numerical simulations to determine $f$ and found \\be \\label{fvalue} f = 10^{4.9 \\pm 0.2} (G \\mu / c^2)^{4.1 \\pm 0.1} \\, . \\ee  The BH are sufficiently small that they lose mass due to the Hawking evaporation process. The fraction of the critical density of the Universe in PBH today due to collapsing CSL is (see \\cite{jru} and references quoted therein) \\be \\label{omegaPBH} \\Omega_{PBH}(t_o) = \\frac{1}{\\rho_{crit}(t_o)} \\int_{t_*}^{t_o} dt' \\frac{dn_{BH}}{dt'} m(t',t_o) \\, , \\ee where $t_o$ is the present age of the Universe; $t_*$ is formation time for a PBH with mass $M_* = 4.4 \\times 10^{14} h^{-0.3} \\; \\mbox{g}$, which is expiring today; $m(t',t_0)$ is the mass of a PBH formed at a time $t'$ at a later time $t$; and $h$ is the Hubble parameter in units of $100 \\mbox{km} \\mbox{s}^{-1} \\mbox{Mpc}^{-1}$.  PBH formed at times $t < t_* \\; (M < M_*)$ do not contribute to this integral because they will have evaporated by today. If we assume for simplicity that PBH with mass $M > M_*$ will have evaporated little by the present time, we can approximate $m(t',t_0)$ by $\\alpha \\mu c t'$. ",
        "conclusions": "We have taken advantage of the recent numerical simulations to better understand PBH formation. The observational consequences of a PBH distribution were used to constrain the CS scenario. We have found that the limits on $G \\mu/c^{2}$ are comparable to those stemming from other criteria. Unless the mass of the BH remnants is larger than $10^3 m_{pl}$, these remnants will contribute negligibly to the dark matter of the Universe, even if the BH formation rate has the maximal value allowed by the $\\gamma$-ray flux constraints. A remnant mass of $10^3 m_{pl}$, however, can arise naturally in some models \\cite{CPW} of BH evaporation. In this case, cosmic strings could consistently provide an explanation for the origin of cosmological structure, for the dark matter, and for the origin of the extragalactic $\\gamma$-ray  and Galactic cosmic ray backgrounds around $100 MeV$."
    },
    "9804/astro-ph9804207_arXiv.txt": {
        "abstract": "We present new elements in the identification of the lens-system producing the 4 images of the BAL quasar H1413+117, based on the recent HST/NICMOS-2/F160W observations. After a careful PSF subtraction of the 4 images of the quasar, the residual H$_{F160W}$ image reveals the presence of a faint object ($H\\sim 20.5$) within the region enclosed by the 4 quasar images. This object corresponds to a single galaxy: the primary lens of the lens-system. We also identify the galaxies around the Cloverleaf which had been proposed to belong to a galaxy cluster/group at high redshift (Kneib et al 1998): the other component in the lens-system that provides the additional ``external\" shear. For these galaxies, we have derived a likely redshift based upon their R$_{F702W}$, I$_{F814W}$ and H$_{F160W}$ magnitudes. We find that most of them are consistent with belonging to a galaxy cluster/group with mean redshift $\\overline{z}=0.9 \\pm 0.1$. Furthermore we detect 2 very red objects ($I-H\\sim 4$): the faintest one has no observed optical (R$_{F702W}$ and I$_{F814W}$) counterpart, while the brightest has a predicted redshift around $z\\sim 2$, and may be identified with one of the Cloverleaf absorbers. This gravitational-lens system constitutes an excellent target for IR imaging/spectroscopy with the new generation of 8m ground-based telescopes. ",
        "introduction": "The excellent quality and broad wavelength coverage of current observations of gravitational lensing systems have allowed to unveil part of their mysteries. The primary lens is often detected. The immediate surrounding of the multiply imaged quasars/galaxies is studied in great detail and generally shows some galaxy clustering ({\\it e.g.} Tonry 1998, Hjorth \\& Kneib 1998) or even sometimes X-ray cluster emissions (Hattori et al 1997, Chartas et al 1998). Furthermore, the measure of a time-delay in Q0957+561/PG1115 (Kundic et al 1997, Schechter et al 1997) has strengthened the interest of a detailed study of these gravitational lens systems in order to use them as a cosmological tool.  Since the identification by Magain et al (1988) of the quadrupole lens-system called the Cloverleaf, 4 images of the BAL quasar H1413+117 at z=2.558, many efforts have been dedicated to a direct search of the lens or of elements of the lens-system. Early models of the lens-system have involved one or two galaxy-lenses very close to the line-of-sight toward the quasar (Kayser et al, 1990). A more recent analysis showed that an external shear was needed to model this system correctly (Keeton, Kochaneck \\& Seljak 1997), and indeed it is probably related to the existence of an overdensity of galaxies nearby, as detected by Kneib et al (1998).  The lensing geometry, amplification and time-delays are quite sensitive to the parameters of the galaxy-lens expected to be located amid the 4 images of the quasar. A positive detection of the galaxy-lens would bring stringent constraints in the modeling of the lens-system.  However, despite relatively deep searches, either in R and I imaging with the HST (Turnshek et al, 1997; Kneib et al, 1998) or K imaging with the Keck telescope (Lawrence et al 1996), the lensing galaxy has not been detected.  The  galaxy cluster recently revealed near the Cloverleaf (Kneib et al, 1998) was assumed to be at $z\\sim 1.7$ as this corresponds to the mean value of the narrow absorption line-systems observed in the quasar spectra (z=1.44, 1.66, 1.87, 2.07 and 2.09: Turnshek et al, 1988; Magain et al, 1988, Monier et al 1998). Combining the IR image with the R$_{F702W}$ and I$_{F814W}$ WFPC-2 images of this system can allow to estimate the likely redshift for the faint galaxies surrounding the Cloverleaf. The recently acquired HST/NICMOS-2/F160W observation consists of a unique dataset to help answer both the existence of the lensing galaxy and to constrain the distance of the nearby galaxies.  The HST/NICMOS-2 data are presented in Sect.2, while in Sect. 3 we discuss the identification of the lensing galaxy after the PSF subtraction of the 4 quasar images. The redshift estimates (derived from photometry) of the surrounding galaxies are explained in Sect.4. The discussion and concluding remarks are provided in Sect.5. Throughout this paper we use $H_0=$50 km/s/Mpc and $\\Omega_0=$1. ",
        "conclusions": "Two main results have been obtained from the NICMOS-2 data. For the first time the galaxy-lens, H1, close to the line of sight toward the quasar, has been identified. Its position with respect to the quasar line-of-sight is found to be similar (within the uncertainties) to the one derived in the various Cloverleaf gravitational lens models ({\\it e.g.} Kneib et al, 1998). It remains difficult to estimate the redshift of the galaxy-lens H1 because the PSF subtraction leaves an increased background noise in the region amid the 4 quasar images. Yet, a redshift estimate around 1.0 or higher is consistent with the H$_{F160W}$ magnitude and the I$_{F814W}$ lower limit magnitude we have derived for H1. Clearly deep spectroscopic data are needed to solve for the determination of its redshift. We find also that there is a unique galaxy-lens, in contradiction to some early models in which two galaxy-lenses had been envisaged (Kayser et al, 1990). Assuming that H1 is around $z\\sim 1$, and has similar colors and absolute magnitude than the nearby galaxies gives a Mass-to-light ratio of M($<5.1$kpc)/L$_B$ $\\sim$ 25 M/L$_{B\\odot}$.  With regard to the Cloverleaf environment, we show that 8 nearby galaxies have a most probable redshift around 0.9, giving credit to the presence of a galaxy cluster/group along the line of sight to the Cloverleaf. In our previous modelling (Kneib et al, 1998), we assumed for this galaxy cluster/group a redshift of 1.7, as a mean of the redshifts of the 4 absorbers silhouetted on the quasar spectrum. This value should be revised. The location of the galaxy-lens is now known from the NICMOS-2 observations and will be implemented in a new model of the lens-system. One of the faint galaxies surrounding the Cloverleaf appears to be at a larger redshift, around 2, and might be related with the absorber at z$=$2.07 or 2.09 (Monier et al 1998). Further IR imaging/spectroscopy of these galaxies should remove the remaining uncertainties of the Cloverleaf lens-system."
    },
    "9804/astro-ph9804031_arXiv.txt": {
        "abstract": "We present data for 18 blazars observed with the X--ray satellite {\\sl ASCA}, half of which were also observed contemporaneously with the EGRET instrument onboard {\\sl Compton Gamma-ray Observatory} as parts of multi-wavelength campaigns.  The observations show a clear difference in the spectra between three subclasses of blazars, namely the High-energy peaked BL Lac objects (HBLs), Low-energy peaked BL Lac objects (LBLs), and quasar-hosted blazars (QHBs). The \\asca X--ray spectra of HBLs are the softest, with the power law energy index $\\alpha \\sim 1 - 2$, and they form the highest observable energy tail of the low energy (LE, synchrotron) component. The X--ray spectra of the QHBs are the hardest ($\\alpha \\sim 0.6$) and are consistent with the lowest observable energy end of the high energy (HE, Compton) component.  For LBLs, the X--ray spectra are intermediate. We find that the radiation process responsible for the HE peak for HBLs {\\sl can} be explained solely by Doppler-boosted Synchrotron-Self-Compton (SSC) emission, with the Doppler factor $\\delta$ consistent with the VLBI and/or $\\gamma$--ray variability data.  For many QHBs, on the other hand, the $\\gamma$--rays {\\sl cannot} be solely due to the SSC mechanism unless $\\delta$ is significantly in excess of that inferred from VLBI data.  We consider an alternative scenario consistent with the measured values of $\\delta$, where the SSC component is still present in QHBs and it dominates in the X--ray band, but it is below the observed $\\gamma$--ray spectrum.  With an assumption that the peak of the SSC emission is on the extrapolation of the X--ray spectrum, and adopting $\\delta$ of 10, we infer the magnetic field $B$ to be 0.1 -- 1 Gauss, and Lorentz factors $\\gamma_{b}$ of electrons radiating at the peak of the $\\nu F(\\nu)$ spectrum of $\\sim 10^{3}$ for QHBs;  this is much lower than $\\gamma_{b} \\sim 10^{5}$ for HBLs, even though the values of $B$ are comparable in the two sub-classes.  This difference of $\\gamma_{b}$ is most likely due to the large photon density expected in QHBs (e.g. from thermal components visible in these objects) as compared with that of HBLs;  Compton upscattering of these photons may well provide the observed GeV flux. ",
        "introduction": "The overall electromagnetic spectra of blazars -- a class of active galactic nuclei (AGNs) that includes BL Lac objects and Optically Violently Variable (OVV) quasars -- are believed to be dominated by Doppler-boosted radiation from relativistic jets pointing closely to our line of sight (\\cite{blandford78}; \\cite{blandford79}; \\cite{urry95} for a review of radio loud AGNs). The VLBI studies of these objects show compact radio cores on milli-arcsecond angular scale with jet-like structures which often show superluminal motion, with apparent speeds $v/c \\sim 5 - 10$ (e.g., \\cite{vermeulen}). Apparent variability time scale and luminosity amplification depend on various powers of the ``beaming'' (Doppler) factor $\\delta$ (e.g., \\cite{lind}), defined via $\\delta  = \\Gamma_{j}^{-1} (1 - \\beta \\cos \\theta)^{-1}$, where $\\Gamma_{j}$ is the bulk Lorentz factor of the emitting matter, $\\beta = v/c$, and $\\theta$ is the angle of motion with respect to the line of sight.  Blazars are commonly detected as $\\gamma$--ray sources. The EGRET instrument onboard the {\\sl Compton Gamma-Ray Observatory} (\\cgro) has so far detected emission in the GeV range from $\\sim$ 50 blazars (\\cite{fichtel94}; \\cite{thompson95}; \\cite{mattox97}; \\cite{mukherjee97}); $\\gamma$--ray emission has been detected up to the TeV range from the nearby BL Lac objects Mkn~421 and Mkn~501 with ground-based Cherenkov telescopes (\\cite{punch92}; \\cite{petry93}; \\cite{quinn96}; \\cite{bradbury97}). As these sources show large luminosity and compact emission regions in the spectral range where the opacity to pair production via $\\gamma\\gamma\\rightarrow e^{+}e^{-}$ is large, it is generally accepted that the $\\gamma$--ray emission is anisotropic and Doppler-boosted as well (\\cite{maraschi92}; \\cite{mattox93}; \\cite{dondi95}; \\cite{buckley96}), suggesting that the {\\sl entire} observed electromagnetic emission arises in the jet. As these objects emit in practically every observable waveband, any study of the structure and physical conditions in the jets requires broad-band spectral observations, which, given the rapid large amplitude flux variability, must be conducted simultaneously.  The overall spectra of blazars have two pronounced components: one peaking at low energies (LE), $10^{13}-10^{17}$ Hz (e.g., \\cite{sambruna96}), and another peaking at high energies (HE), in the $\\gamma$--rays (e.g., \\cite{montigny95}). For the blazars that are hosted in quasars (QHBs), and for BL Lac objects discovered via radio-selection techniques (the so-called ``Low-energy peaked BL Lacs'' or LBLs), the LE component peaks in the infrared. For the the majority of BL Lac objects -- those found as a result of their X--ray emission -- it peaks in the ultraviolet or even in the soft X--rays (\\cite{giommi95}; \\cite{sambruna96}; \\cite{padovani96}; \\cite{fossati97a}), and thus they are named ``High-energy peaked BL Lacs (HBLs)'' (\\cite{padovani95}). The local power-law shape, the smooth connection of the entire radio - to - UV (and, for the HBLs, soft X--ray) continuum, as well as the relatively high level of polarization observed from radio to the UV, imply that the emission from the LE component is most likely produced via the synchrotron process of relativistic particles radiating in magnetic field. This view is strongly supported by spectral variability observed in X--rays in a number of HBLs, where the variability at lower X--ray energies lags behind the more energetic X--rays (\\cite{tashiro92}; \\cite{sembay93}; \\cite{kohmura94}; \\cite{tashiro95}; \\cite{takahashi96})  The HE component, on the other hand, peaks in the $\\gamma$--ray band, in the MeV - to - GeV range, and, in the case of a few HBLs, it extends to the TeV range;  it is believed to be produced via Comptonization by the same particles that radiate the LE component. The source of the ``seed'' photons, can either be the synchrotron radiation, internal to the jet -- as in the Synchrotron-Self-Compton (SSC) models (\\cite{rees67}; \\cite{jones74}; \\cite{konigl81}; \\cite{ghisellini85}; \\cite{band85}; \\cite{ghisellini89}; \\cite{maraschi92}; \\cite{bloom96}; \\cite{mastichiadis97}). Alternatively, these can be external to the jet, as in the External Radiation Compton (ERC) models: either the UV accretion disk photons (\\cite{dermer92}; \\cite{dermer97}), or these UV photons reprocessed by the emission line clouds and/or intercloud medium (\\cite{sikora94}; \\cite{blandford95}), or else, IR radiation ambient to the host galaxy (\\cite{sikora94}). The ratio of the power in the HE to the LE components is systematically larger for QHBs than for BL Lac objects (\\cite{maraschi94a}; \\cite{dondi95}; \\cite{sambruna97a}; \\cite{ulrich97}; \\cite{fossati97b}).  If we assume that the LE component is due to the synchrotron radiation, its peak frequency is determined by the intensity of magnetic field and the distribution function of electron energies, while the location of the HE peak is related to the distribution functions of electron and target photon energies. The ratio of the luminosity of these components ($L_{HE}/L_{LE}$), in the context of this synchrotron plus Compton model, is expected to reflect the ratio of energy densities of photon and magnetic fields in the jet.  This paper reports the X--ray spectra of 18 blazars measured by \\asca in the context of their multi-band emission.  The \\asca observations and results are described in \\S2, followed by multi-band analysis and discussion in \\S3. Summary of this paper is presented in \\S4. Throughout this paper we use $H_0$=75 km s$^{-1}$ Mpc$^{-1}$, $q_0$=0.5. ",
        "conclusions": "As we mentioned previously, the two leading models of the high energy emission in blazars invoke Comptonization, of internal (SSC) or external (ERC) seed photons.  In the following analysis, we assume that {\\sl both} SSC and ERC processes may operate in blazars.  We then estimate the contribution of the SSC emission in the HE component.  In order to calculate the predicted luminosity due to the SSC emission, we assume a simple homogeneous model, in which photons are produced in a region of radius $R$ and with a constant magnetic field $B$.  We considered the radiation by a single population of relativistic electrons, with a broken power law distribution of Lorentz factors $\\gamma_{el}$, and a break point at $\\gamma_{b}$ (similar to e.g. \\cite{sambruna96}).  We also assume that the radiation spectrum of the LE component peaks at a frequency corresponding to that radiated by the electrons with $\\gamma_{b}$.  The peak frequency of the synchrotron component in the observer frame, $\\nu_{sync}$, is then given as, when pitch angle is $\\pi$/2: \\begin{equation} \\nu_{sync} = 1.2 \\times 10^{6} \\gamma_{b}^{2} B \\frac{\\delta}{(1+z)} \\quad {\\rm Hz} \\label{eqn1} \\end{equation}  where B is in Gauss.  If the electron energy is still in the Thomson regime, ($\\gamma_{el} \\times h\\nu_{sync} << m_ec^2$), the expected peak of the SSC component in the observer frame ($\\nu_{SSC}$) is $\\nu_{SSC} = 4 \\gamma_{b}^{2} \\nu_{sync}/3$. The ratio of the observed luminosity of the SSC component $L_{SSC}$ to the observed synchrotron luminosity $L_{sync}$ is: \\begin{equation} \\frac{L_{SSC}}{L_{sync}}=\\frac{u_{sync}}{u_{B}} \\label{eqn2} \\end{equation} where the $u_{sync}=L_{sync}/(4\\pi R^{2}c\\delta^{4})$ is the rest-frame energy density of the synchrotron photons, and $u_{B}=B^{2}/(8\\pi)$ is the magnetic field energy density.  To check the validity of the assumption that the observed HE component is solely due to the SSC emission, we calculated the beaming factor ($\\delta$), which is given from above equations: \\begin{equation} \\delta^2 = 1.6\\times 10^{12} \\frac{L_{sync}}{c R^2} \\left(\\frac{L_{sync}}{L_{SSC}}\\right) \\frac{\\nu_{SSC}^2}{\\nu_{sync}^4} \\frac{1}{(1+z)^2} \\label{eqn3} \\end{equation}  where $L$ is in erg~s$^{-1}$, $\\nu$ in Hz, $c$ in cm s$^{-1}$, and $R$ in cm. We estimate $R$ from the shortest observed variability (doubling) time scale $\\Delta t$ observed in any wavelength, as given in Table 2.  Assuming that $R \\lesssim c \\delta \\Delta t / (1+z)$, then Eq. 3 can be rewritten as: \\begin{equation} \\delta^4 \\gtrsim 1.6\\times 10^{12} \\frac{L_{sync}}{c^3 \\Delta t^2} \\left(\\frac{L_{sync}}{L_{SSC}}\\right) \\frac{\\nu_{SSC}^2}{\\nu_{sync}^4} \\label{eqn4} \\end{equation}  where $\\Delta t$ is in s, and other quantities are as in Eq. 3. An application of this equation to the data in Table 2 assuming $L_{sync}=L_{LE}$, $L_{SSC}=L_{HE}$, $\\nu_{sync}=\\nu_{LE}$, $\\nu_{SSC}=\\nu_{HE}$ implies that the lower limits of $\\delta$ for HBLs are $\\sim $3 or less, which is consistent with the VLBI results (cf. \\cite{vermeulen}), and the limits obtained from the arguments of the $\\gamma$--ray opacity (cf. \\cite{dondi95}).  However, for 4 QHBs, where the $\\gamma$--ray flux severely dominates the radiative output, we derive values of $\\delta$ that are much larger than the VLBI results (see Table 2).  This suggests that an additional emission mechanism -- such as the ERC process -- may contribute significantly in the $\\gamma$--ray regime, dominating over the SSC flux, and the values of $\\nu_{SSC}$ and $L_{SSC}$ are {\\sl very} different than $\\nu_{HE}$ and $L_{HE}$, with the SSC component ``hidden'' well below the ERC component.  However, the fact that the QHBs have X--ray spectra which are hard, with $\\alpha \\sim 0.6$, and which are {\\sl not} located on the extrapolation of the synchrotron optical / UV spectra, implies that the X--rays observed in QHBs are due to a separate emission process than synchrotron. The fact that for most of QHBs the $\\gamma$--ray spectra are above the extrapolation of X--ray spectra (Fig. 2) suggests that the dominant process is different for X--rays than it is for $\\gamma$--rays. One explanation is that the SSC process dominates in the X--ray range, while the ERC process dominates in $\\gamma$--rays (\\cite{inoue}). With the assumption that SSC process is dominant in X--rays for QHBs, we estimate the location of the ($\\nu_{SSC}$, $L_{SSC}$) point in the log($\\nu$) -- log($\\nu F (\\nu)$) space by the following method. We assume that it lays on or below the extrapolation of the {\\sl ASCA} spectrum (line (a) in Fig. 3), but above the highest value of $\\nu F (\\nu)$ measured by {\\sl ASCA}. Since the spectra of QHBs generally have $\\alpha < 1$ and thus $\\nu F (\\nu)$ is the highest at the end of the \\asca bandpass (10 keV $\\simeq$ 2$\\times$10$^{18}$Hz), this second limit is equivalent to $\\nu_{SSC} L_{SSC} > 2\\times 10^{18}$ Hz $L_{10 \\rm keV}$ (line (b) in Fig. 3). We further constrain $L_{SSC}$ using Eq. 3; once we assume a given $\\delta$, there is a unique relationship between $L_{SSC}$ and $\\nu_{SSC}$ described as: \\begin{equation} L_{SSC} = 1.6 \\times 10^{12} \\left(\\frac{L_{sync}^{2}}{cR^{2}\\nu_{sync}^{4} \\delta^{2} (1+z)^2} \\right) \\nu_{SSC}^{2} \\quad {\\rm erg~s^{-1}} \\label{eqn5} \\end{equation} where $L$, $R$, $c$, and $\\nu$ are in the same units as in Eqs 3 \\& 4. The VLBI data and $\\gamma$-ray opacity argument suggest that $5 < \\delta < 20$ for most blazars (e.g.,\\cite{vermeulen}; \\cite{dondi95}). This corresponds to the lines (c) and (d) in Fig. 3, respectively for $\\delta$=5 and 20. The above four constraints correspond to the shaded area of Fig. 3, where for illustration, we use the overall spectral energy distribution for the QHB CTA~102.  Since we have to use a unique value in calculating the physical parameters, we use $\\delta$ = 10 as a geometrical mean between 5 and 20.  The $L_{SSC}$ - $\\nu_{SSC}$ line calculated from Eq. 5 corresponding to $\\delta$ = 10 intersects both the extrapolation of the {\\sl ASCA} spectrum and the highest {\\sl ASCA} value, and the intersections yield the lower and upper values for both $L_{SSC}$ and $\\nu_{SSC}$.  We adopt a mean of these values, which are given in Table 2, and plotted in Figure 4c. We used \\ginga data instead of \\asca data for 3C~279 because a simultaneous campaign from radio to $\\gamma$--ray bands was conducted during \\ginga observation. The values for the other blazars where $\\delta$ derived from Eq. \\ref{eqn4} is $<20$, are calculated by assuming $L_{SSC}$ = $L_{HE}$. For LBL AO0235+164 where $\\delta$ derived from Eq. 4 is $>20$, we assume $L_{SSC}$ = $L_{LE}$ because the \\asca spectra of AO0235+164 is thought to be mixture of the LE and HE component, as discussed by \\cite{madejski96} based on the \\rosat and \\asca spectra, so that the above method may be inappropriate. For two QHBs (3C~273, PKS~0208-512) the lower limits of $\\delta$ are $\\sim$5. The fact that observed $\\gamma$-ray flux of PKS~0208-512 is much higher than the extrapolation of \\asca spectrum implies the $\\gamma$-ray peak is not solely due to SSC mechanism. Therefore we applied the above method to this source. On the other hand, since the $\\gamma$-ray spectrum of 3C~273 is below the extrapolation of \\asca spectrum, the $\\gamma$-ray emission is assumed to be due to SSC mechanism so there, we assume $L_{SSC}$ = $L_{HE}$.  It is important to note, however, that 3C~273 is unique as compared to other blazars considered here in that the ``blue bump'' is very pronounced, and thus it is unlikely that the jet dominates the entire electromagnetic emission, and therefore, a more complex analysis is necessary (see, e.g., \\cite{montigny97} for further discussion).  Once we obtain $L_{SSC}$, and $\\nu_{SSC}$, we can calculate the strength of the magnetic field and the electron Lorentz factor $\\gamma_{b}$ from Eq. (\\ref{eqn1}), (\\ref{eqn2}) and those are given as follows: \\begin{equation} B = 0.27 \\left(\\frac{R_{\\mbox{pc}}}{10^{-2}}\\right)^{-1} \\left(\\frac{\\delta}{10}\\right)^{-2} \\sqrt{\\left(\\frac{L_{sync}}{10^{46}}\\right) \\left(\\frac{L_{sync}}{L_{SSC}}\\right)} \\quad {\\rm Gauss} \\end{equation} \\begin{equation} \\gamma_{b} = 1.8\\times 10^3 \\left(\\frac{R_{\\mbox{pc}}}{10^{-2}}\\right)^{1/2} \\left(\\frac{\\delta}{10}\\right)^{1/2} \\left(\\frac{\\nu_{sync}(1+z)}{10^{13}}\\right)^{1/2} \\left[\\left(\\frac{L_{sync}}{10^{46}}\\right) \\left(\\frac{L_{sync}}{L_{ssc}}\\right)\\right]^{-1/4} \\end{equation} where $R_{\\mbox{pc}}$ is size of emission region in parsecs, and other quantities are as in Eqs 3, 4, \\& 5. As before, the upper limit of the size $R$ can be estimated from the observed time variability ($\\Delta t$) from Table 2, given by $R \\lesssim c\\Delta t \\delta / (1+z)$.  Our calculated values of $B$ and $\\gamma_{b}$ are plotted respectively in Figures 4d and 4e.  In these Figures, we also plot the values calculated with $R = 0.01$ pc, which would correspond to an observed variability time scale of $\\sim 1$ day.  From our analysis, the magnetic field for blazars observed with \\asca is inferred to be 0.1 -- 1 Gauss. The value of $B$ is comparable between the different subclasses of blazars, although $B$ is somewhat lower in HBLs than in QHBs. With these values of $B$, we estimate $\\gamma_{b}$ to be $10^{3} - 10^{4}$ for QHBs, and $10^{5}$ for HBLs.  The differences of $\\gamma_{b}$  between different sub-classes of blazars imply that the relativistic electrons are accelerated to higher energies in HBLs than in QHBs. Alternatively, higher $\\gamma_{b}$ in HBLs might be obtained by increasing $\\delta$.  However, in those objects, we believe that there is no contribution to the $\\gamma$--ray production from other mechanisms besides SSC, and thus the observed $L_{HE}$ is $L_{SSC}$.  In such case, $\\gamma_{b}$ depends on $\\delta$ only linearily (cf. Eq. 7 and $R\\lesssim c\\Delta t \\delta/(1+z)$), and thus varying $\\delta$ to be 5 or 20 respectively decreases or increases our derived $\\gamma_{b}$ only by a factor of two, which is small when compared to the large difference of $\\gamma_{b}$ calculated by us (cf. Fig. 4e).  In QHBs the strong optical and UV line emission implies a presence of dense external radiation fields.  This means that in the frame of reference of the jet, these can easily dominate over the internal synchrotron radiation, resulting in the ERC emission dominating over the SSC emission in $\\gamma$--ray band (e.g., \\cite{sikora97}). It is likely that the difference of $\\gamma_{b}$ is most likely due to the large photon density in QHBs as compared with that of HBLs. It should be noted that TeV photons have been observed only from HBLs, where we calculate higher values of $\\gamma_{b}$."
    },
    "9804/astro-ph9804096_arXiv.txt": {
        "abstract": "In this paper we present a new method that can be used for analysis of time of arrival  of a pulsar pulses (TOAs). It is designated especially to detect quasi-periodic variations of TOAs. We apply our method to timing observations of PSR B1257+12 and demonstrate  that using it it is possible to detect  not only first harmonics of  a periodic variations, but also  the presence of a resonance effect. The resonance effect detected, independently of its  physical origin, can appear only when there is a non-linear interaction between two periodic modes. The explanation of TOAs variations as an effect of the existence of planets is, till now, the only known and well justified. In this context, the existence of the resonance frequency in TOAs is the most significant signature of the gravitational interaction of planets. ",
        "introduction": "The first extra-solar planetary system was discovered by \\cite{Wolszczan:92::} around a millisecond radio pulsar, PSR B1257+12. The three planets orbiting the pulsar have been indirectly deduced from the analysis of quasi-periodic changes in the times of arrival (TOAs) of pulses caused by the pulsar's reflex motion around the center of mass of the system. In the analyses of this kind, it is particularly important to establish a reliable method of distinguishing planetary signatures from possible TOA variations of physically different origin. In the case of PSR B1257+12, it was possible to make this distinction and confirm the pulsar planets through the detection of mutual gravitational perturbations between planets B and C \\cite[]{Wolszczan:94::}, following predictions of the existence of this effect by \\cite{Rasio:92::}, \\cite{Malhotra:92::} (see also \\cite{Malhotra:93::}, \\cite{Rasio:93::} and \\cite{Peale:93::}).  Practical methods of detection of the TOA variations caused by orbiting planets include direct fits of Keplerian orbits to the TOA or TOA residual data \\cite[e.g.]{Thorsett:92::,Lazio:95::} and model--independent frequency domain approaches based on Fourier transform techniques \\cite[]{Konacki:96::,Bell:97::}. In fact, it appears that it is best to search for periodicities in TOAs (or post-fit TOA residuals left over from fits of the standard timing models) by examining periodograms of the data, and then refine the search by fitting orbits in the time domain using initial orbital parameters derived from a frequency domain analysis.  The presence of planets around a pulsar causes pulse TOA variations which, for planets moving in orbits with small eccentricities, have a quasi-periodic character and generate predictable, orbital element-dependent features in the spectra of TOA residuals. This has led \\cite[]{Konacki:96::} to devising a method of TOA residual analysis based on the idea of a successive elimination of periodic terms applied by Laskar (1992) in his frequency analysis of chaos in dynamical systems. The frequency analysis provides an efficient way to decompose a signal representing the TOA residual variations into its harmonic components and study them in an entirely model-independent manner. As shown in \\cite[]{Konacki:96::} this method works perfectly well under idealized conditions in which covariances among different parameters of the timing model are negligible.  In this paper, we present an improved scheme for the frequency analysis of pulsar timing observations in which a successive elimination of periodicities in TOAs is incorporated in the modelling process rather than being applied to the post-fit residuals. This makes the results obtained with our method less sensitive to the effect of significant covariances which may exist between various timing model parameters. We apply the frequency analysis to TOA measurements of the planet pulsar, PSR B1257+12, using a computer code developed to fit spectral timing models to data. We show that our method allows an easy detection of the fundamental orbital frequencies of planets A, B and C in the pulsar system and the first harmonics of the frequencies of planets B and C generated by eccentricities of the planetary orbits. Furthermore, by detecting the effect of perturbations between planets B and C, we demonstrate that the frequency analysis method represents a sensitive, model-independent tool to analyze nonlinear interactions between periodic modes of processes of various physical origins. ",
        "conclusions": ""
    },
    "9804/astro-ph9804269_arXiv.txt": {
        "abstract": "We have developed in detail the theory of X-ray line and continuum production due to atomic interactions of accelerated ions, incorporating in our calculations information from a broad range of laboratory measurements. We applied our calculations to the Orion region from which nuclear gamma-ray lines were observed with the COMPTEL instrument on {\\it CGRO}. The accelerated particles which produce this gamma-ray emission via nuclear reactions also produce X-ray lines via atomic interactions. We predict strong line emission in the range from 0.5 to 1 keV, mainly due to de-excitations in fast O ions. While much of the diffuse X-ray emission observed with ROSAT from Orion could be due to accelerated ions, the current X-ray data do not provide unambiguous signatures for such an origin. If future observations with high spectral resolution would confirm the predicted X-rays, the combined analysis of the X-ray and gamma-ray data will set important constraints on the origin of the accelerated particles and their interaction model. ",
        "introduction": "Strong gamma-ray emission in the 3-7 MeV range has been detected from the Orion molecular cloud complex with the COMPTEL instrument on the {\\it Compton Gamma Ray Observatory (CGRO}; Bloemen et al. 1994, 1997). As the observed spectrum exhibits characteristic structures (Bloemen et al. 1997), this emission is most likely due to the superposition of nuclear gamma-ray lines, mainly the 4.44 MeV line from $^{12}$C and the 6.13, 6.92 and 7.12 MeV lines from $^{16}$O. Such line emission can only be produced by accelerated particle interactions. Gamma-ray emission at photon energies $>$30 MeV was also observed from Orion, with the EGRET instrument on {\\it CGRO} (Digel, Hunter, \\& Mukherjee 1995). This gamma-ray emission is consistent with pion production and bremsstrahlung due to irradiation by standard Galactic cosmic rays (Digel et al. 1995). As such cosmic rays underproduce the observed line emission by at least three orders of magnitude, the gamma-ray line production in Orion must predominantly be a low energy cosmic ray phenomenon. Information on the spatial distribution of the gamma-ray line emission in Orion has come from both the COMPTEL and {\\it CGRO}/OSSE observations. The extended nature of the emission seen in the COMPTEL map of Orion (Bloemen et al. 1997) could provide an explanation for the fact that so far it was not possible to confirm the COMPTEL results with OSSE (Murphy et al. 1996; Harris et al. 1998).  Based on the observed line widths, Bloemen et al. (1994) first suggested that the line emission is produced by accelerated C and O ions interacting with ambient H and He, rather than by accelerated protons and $\\alpha$-particles interacting with ambient C and O. More detailed analyses of the initial COMPTEL data have shown that a mix of the two processes could not be ruled out (Ramaty, Kozlovsky, \\& Lingenfelter 1995; Cowsik \\& Friedlander 1995). But, as the emission peaks in the more recent COMPTEL data do not appear at the line center energies for $^{12}$C and $^{16}$O de-excitations (Bloemen et al. 1997), a significant narrow-line contribution from accelerated proton and $\\alpha$-particle interactions seems to be excluded (Kozlovsky, Ramaty, \\& Lingenfelter 1997). This conclusion is also supported by energetic arguments, as the very large power deposited by the accelerated particles into the ambient medium in Orion is lowered by enhancing the C-to-proton and O-to proton abundance ratios (Ramaty et al. 1995; Ramaty, Kozlovsky \\& Lingenfelter 1996). Apart from the observed emission in the 3-7 MeV band, the COMPTEL observations revealed only upper limits at other gamma-ray energies (Bloemen et al. 1994, 1997). In particular, the upper limit on the 1-3 MeV emission sets constraints on the accelerated Ne-Fe abundances relative to those of C and O. The suppression of both the Ne-Fe and proton and $\\alpha$-particle abundances relative to C and O could be understood if the seed particles injected into an as-yet unknown particle accelerator (see Nath \\& Biermann 1994; Bykov \\& Bloemen 1994) come from the winds of massive stars or the ejecta of supernovae resulting from massive progenitors (Bykov \\& Bloemen 1994; Ramaty et al. 1995; Cass\\'e, Lehoucq, \\& Vangioni-Flam 1995; Ramaty et al. 1996; Parizot, Cass\\'e, \\& Vangioni-Flam 1997a). Ip (1995) and Ramaty et al. (1996) have also considered the possible acceleration of ions resulting from the breakup of interstellar dust.  The gamma-ray line production in Orion should be accompanied by a large ionization rate of the ambient medium which could exceed the observed infrared luminosity (Cowsik \\& Friedlander 1995). This problem is alleviated if the gamma-rays are produced at cloud boundaries, but not in their interiors. The accelerated particles could have ionized $\\sim$2$\\times$10$^4$M$_\\odot$ in 10$^5$ years (Ramaty 1996), a small fraction of the total available mass. It is thus possible that a large fraction of the power that accompanies the gamma-ray production is deposited in an ionized gas.  While the X-ray emission produced by low energy particle interactions is potentially a promising tracer of low energy cosmic rays in the Galaxy (e.g. Hayakawa \\& Matsuoka 1964), there are as-yet no astrophysical X-ray observations that unambiguously indicate the presence of such cosmic rays. The Orion region, however, has become an interesting target owing to the COMPTEL discovery of the nuclear gamma-ray line emission. A variety of processes lead to X-ray production by low energy ion interactions. Inverse bremsstrahlung (Boldt \\& Serlemitsos 1969) results from the interactions of fast ions and ambient electrons; secondary electron bremsstrahlung is produced by knock-on electrons accelerated in fast ion interactions (Hayakawa \\& Matsuoka 1964). Both of these processes lead to continuum X-ray emission. X-ray line emission results from atomic de-excitations in the fast ions following electron capture (Silk \\& Steigman 1969; Watson 1976; Pravdo \\& Boldt 1975; Bussard, Ramaty, \\& Omidvar 1978) and in ambient ions following inner-shell vacancy creation. The latter process has not yet been applied to astrophysics. Dogiel et al. (1997) have recently considered the X-ray emission that should accompany the gamma-ray line production in Orion. They have only considered the secondary electron bremsstrahlung and concluded that the 0.5-2 keV emission that accompanies the observed gamma-ray line emission from Orion will exceed the upper limits that they derived using ROSAT observations. We have subsequently taken into account both continuum processes and line emission from de-excitations in fast O (Ramaty, Kozlovsky, \\& Tatischeff 1997a) and showed that, even though the inverse bremsstrahlung is more important than the secondary electron bremsstrahlung, the total X-ray continuum emission from Orion is not inconsistent with the Dogiel et al. (1997) derived ROSAT upper limit. On the other hand, we showed that a conflict may exist between that ROSAT upper limit and the X-ray line emission following electron capture onto fast O nuclei. However, as we suggested, this conflict could be resolved if the X-ray and gamma-ray lines are produced in an ionized medium or if the current epoch accelerated particle spectrum is suppressed at low energies, for example by energy losses.  In this paper we present detailed calculations of X-ray continuum and line production by accelerated particle interactions. The bulk of our treatment is for a steady state, thick target model with a neutral ambient medium. This is the standard model in which most of the gamma-ray calculations have been carried out (e.g. Ramaty et al. 1996). But we have also investigated the effects of an ionized ambient medium and a time-dependent model, as these modifications could have important consequences on the predicted X-ray to gamma-ray production ratio. In our treatment of the continuum, we have supplied the details of the calculations and we have improved the employed cross sections, thereby confirming our previous preliminary results (Ramaty et al. 1997a). We have greatly expanded our treatment of X-ray line emission. We have investigated in detail the atomic physics relevant to line emission from de-excitations in fast O, checking our theoretical calculations against laboratory data whenever available. We then expanded the treatment to the other abundant accelerated ions (C, N, Ne, Mg, Si, S and Fe), and we have also calculated the X-ray line emission produced in ambient ions following inner-shell vacancy creation by the accelerated particles. We have used the ROSAT all-sky survey (Snowden et al. 1995) to derive the X-ray count rates from the Orion region that could be associated with accelerated particle interactions; the implied fluxes are quite different from the upper limit given by Dogiel et al. (1997). The unambiguous future detection of the predicted X-rays produced by accelerated particles in Orion, and potentially elsewhere in the Galaxy, should provide important new insights into the origin of the low energy cosmic rays whose presence in Orion is revealed by the COMPTEL gamma-ray line observations. ",
        "conclusions": "We have investigated all the processes that lead to X-ray production by low energy cosmic rays for a variety of accelerated particle compositions and energy spectra. We demonstrated that the dominant continuum producing process is inverse bremsstrahlung produced by fast ions interacting with ambient electrons. In addition, there is also a significant contribution from the bremsstrahlung produced by secondary knock-on electrons. However, below a few keV the total X-ray emission produced by accelerated ions is dominated by relatively broad line emission (line widths $\\delta E/E$$\\simeq$0.1) resulting from de-excitations in the fast ions following electron captures and excitations. In addition, accelerated particle interactions also produce much narrower X-ray lines, due to inner-shell vacancy creation. The most prominent of such line is that at 6.4 keV from ambient Fe.  We have calculated the X-ray line and continuum emission produced by the accelerated particles in Orion which are thought to be responsible for the nuclear gamma-ray line emission observed with COMPTEL (Bloemen et al. 1994, 1997). By first comparing the results with the extragalactic diffuse X-ray background, we found that while the continuum is generally below this background, the line emission from about 0.5 to 1.5 keV exceeds the background for all the combination of parameters that we considered. We wish to point out that there could be a significant contribution to the $\\sim$0.5-1.5 keV diffuse X-ray background from as-yet unknown sources within our Galaxy (e.g. Park et al. 1997), leaving the possibility that a substantial fraction of the observed X-ray intensity in this energy range results from low energy cosmic ray interactions.  We have also compared our results with ROSAT observations of Orion in the 0.47 to 1.2 keV energy band, again normalizing the X-ray emission to the observed gamma-ray emission. We found that there is no conflict between the predicted total X-ray emission (lines and continuum) and the data for a broad range of parameters if the gamma-ray line emission is uniformly distributed over the entire molecular cloud complex. This conclusion differs from our previous one (Ramaty et al. 1997a) because of a lower predicted X-ray line emission, resulting from improved atomic physics input, and because our estimated ROSAT flux from Orion is higher than the upper limit given by Dogiel et al. (1997). However, the COMPTEL data show significant spatial structure. We found that for the most prominent hot spot in the COMPTEL map, the standard thick target, steady state interaction model, with a neutral ambient medium, predicts X-ray fluxes which exceed the ROSAT data for a broad range of parameters. But the calculations could be consistent with the data for any one, or a combination of the following possibilities: a very hard accelerated particle spectrum; a partially ionized ambient medium; and a time-dependent accelerated particle energy spectrum resulting from  essentially instantaneous acceleration some tens of thousand of years ago.  There are as-yet no astrophysical X-ray observations that would unambiguously indicate an origin resulting from low energy, accelerated ion interactions. Our calculations show that the most promising signatures are the relatively broad lines between 0.5 and 1.5 keV, mainly the lines from fast O, and that a promising target is the Orion region where the presence of such accelerated particles is known from gamma-ray line observations.  We acknowledge K. Omidvar for discussions on the atomic processes leading to X-ray line production. V. T. acknowledges an NRC-NASA/GSFC Research Associateship.  \\clearpage"
    },
    "9804/astro-ph9804025_arXiv.txt": {
        "abstract": "We report the serendipitous discovery of a 7-s X-ray pulsar using data acquired with the {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca ). The pulsar is detected as an unresolved source located towards a region of the Galactic plane ($l,b \\simeq 29.5, 0.08$) that coincides with an overdensity of star-formation tracers. The signal suffers tremendous foreground absorption, equivalent to $N_H \\simeq 10^{23}$ cm$^{-2}$; the absorption correlates well with a line-of-sight that is tangential to the inner spiral arms and the 4-kpc molecular ring. The pulsar is not associated with any known supernova remnants or other cataloged objects in that direction. The near sinusoidal pulse (period $P \\simeq 6.9712$) is modulated at 35\\% pulsed amplitude, and the steep spectrum is characteristic of hot black-body emission with temperature $kT \\sim 0.65$ keV. We characterize the source as an anomalous X-ray pulsar (AXP). ",
        "introduction": "A canonical young pulsar (period $\\sim 100$ ms and stellar dipole field $\\sim 3\\times10^{12}$ G) is a rapidly rotating neutron star, created as the stellar remnant during a Type II (or Ib) supernova explosion of a massive star. The birthrate of pulsars is known to be close to $1 - 3$ per century, and it is estimated that there are about $10^5$ active and $10^8$ defunct neutron stars in the Galaxy (see Lorimer et al. 1993 and refs. therein).  In the last few years, there has been growing recognition of a population of ultra-magnetized neutron stars, or ``magnetars'' (Thompson \\& Duncan 1993).  The mostly circumstantial evidence comes from investigations of the following categories of objects: the soft gamma-ray repeaters (Thompson \\& Duncan 1995; Frail et al. 1997), long period pulsars in supernova remnants (Vasisht \\& Gotthelf 1997 and refs. therein), other seemingly isolated, young, long period pulsars (Thompson \\& Duncan 1996) nowadays referred to as the anomalous X-ray pulsars (AXP; van Paradijs et al. 1995), and perhaps their older variants (Kulkarni \\& van Kerkwijk 1998). These objects share some common properties; they are steady, bright X-ray sources ($L_X \\simgt 10^{35}$ erg s$^{-1}$) which show no signs for an accompanying companion, those with known periods are found to be spinning down, and all are relatively young ($ \\simlt 10^{5}$ yr-old).  The evolutionary consequences of such large dipole fields are reflected in the properties listed above.  Most importantly, large braking torques acting on the star cause it to spin-down rapidly, and the magnetic free energy quickly dominates over the rotation energy, i.e., within several hundred years. For the above sources, the rotation rates lie between 6 - 12 s, with ages $\\simlt 10^5$ yr (for the SGRs the evidence for periods is indirect, however, their ages are well constrained due to their association with supernova remnants). It is believed that field decay, which is expected for ultramagnetized neutron stars (Thompson \\& Duncan 1996; also Goldreich \\& Reisenegger 1992), influences the thermal evolution and powers the large X-ray luminosities observed for the purported magnetars, $L_X \\simgt 10^{35}$ erg s$^{-1}$.  If magnetars represent the tail-end of the magnetic field distribution of neutron stars, then they are bound to be rare. Assume that their birthrate is 10\\% the birthrate of neutron stars (some justification for this comes from the estimated birthrates of SGRs; Kulkarni et al. 1994), and that they have active X-ray lifetimes of $\\sim 10^{5}$ yr. These assumptions imply that at present there are only $\\sim 100$ active magnetars in the Galaxy, a conclusion that is borne out by the observations of the aforementioned objects. The fact that we observe the five known AXPs through large column densities in the Galaxy, $N_H \\simgt 10^{22}$ cm$^{-2}$, suggests that they are indeed that rare, and the fact that they are often associated with supernova remnants or lie near star-formation regions (in spite of the large random velocities usually attributed to neutron stars; Lyne \\& Lorimer 1994) suggests that they are young. Similarly, only two Galactic SGRs are known, and it has been suggested that the SGR population census is nearly complete (Kouveliotou 1995).  In summary, AXPs have long rotation periods, hot blackbody-like spectra ($kT \\sim$ 0.5 keV) with $10^{35-36}$ erg s$^{-1}$ steady luminosities, and have thus far only been observed at X-ray wavelengths. A search for new AXPs in the ASCA database has turned up another candidate, which we refer to as \\psr. ",
        "conclusions": "On the basis of its long rotation period, steady X-ray flux, steep spectral characteristics, and location in the Galactic plane ($|b| \\le 0.5$), we classify \\psr\\ as an anomalous X-ray pulsar (see Table 1). The high foreground absorption suggests that the pulsar is distant, and its line of sight along the tangent to the Sagittarius-Carina and Scutum-Crux spiral arms, and the 4-kpc molecular ring justifies its enormous foreground absorption (see figure 4). Its $N_H$ is roughly twice that of the nearby remnant Kes 73 for which quoted distances lie between $10 - 20$ kpc (Blanton \\& Helfand 1996). However, the disparity in $N_H$ does not in itself imply a great dissimilarity in distances. For instance, at 10 kpc the lines of sight vectors to these two objects are already separated by $\\sim 100$ pc, the typical sizes and scale heights of dense giant molecular clouds which are likely to be responsible for most of the absorbing gas. For the purposes of this article we assume the distance to be 15 kpc, an estimate likely to be accurate to within a factor of two.  The steep X-ray spectrum is characteristic of the Wien tail of a blackbody radiator. Using the best fit blackbody parameters the isotropic X-ray luminosity is $L_X \\simeq 2.5\\times 10^{35}d_{15}^2$ erg s$^{-1}$, the distance being 15$d_{15}$ kpc. In effect, the X-ray pulsations can be ascribed to the viewing of a rotating stellar hotspot of area $0.15 A_sd_{15}^2$, where $A_s$ is the area of a neutron star of radius 10 km; this estimate ignores any relativistic corrections to the inferred area. Note that the spectrum is unlike that of any accreting high-mass neutron star binary. Although such binaries have periods in the range 0.07 - 900 s, and sometimes go into low luminosity states with $L_X \\sim 10^{35}$ erg s$^{-1}$, they generally display very hard spectra ($0.8 < \\Gamma < 1.5$), and show stochastic variability on all time-scales (Nagase 1989; Koyama et al. 1989) as is generally seen in accretion powered sources. We find no evidence for such variability in our data.  With an AXP classification in hand we can compare the properties of \\psr\\ with those of five other members of the AXP family in Table~1. A few years ago Schwentker (1994) reported weak 5-s pulsations from RX J1838.4$-$0301 which have not yet been confirmed, Mereghetti et al. (1997) have argued that this X-ray source might be due to coronal emission from a late type star. We, therefore, exclude this source from our list. Although only a future $\\dot P$ measurement can help determine the linear spindown age of \\psr\\ (an estimator for the age of an isolated neutron star), its location in the Galactic plane suggests that it is young, $\\tau < 10^5$ yr. We consider it extremely likely that the pulsar is associated with one of the several star-formation complexes expected to lie along this line-of-sight (see Fig 4), out to a distance of 20 kpc. The pulsar lies along a rich region of the Galaxy; there are 10 supernova remnants, several radio pulsars, and the $\\gamma$-ray source GRO J1838$-$04, all within a $3\\times3$ deg$^2$ patch of sky surrounding \\psr. However, to the best of our knowledge no cataloged sources are associated with it.  We now describe other models that address the unique properties of AXPs. It is often noted in the literature, that the inferred accretion rate for the pulsars in Table~1 are close to those expected for accretion powered pulsars with field strengths $B \\sim 10^{11-12}$~G, spinning at their equilibrium periods $P_{eq}$ (see Bhattacharya \\& van den Heuvel 1991). This motivated Mereghetti \\& Stella (1995) to suggest that these pulsars are members of a subclass of low mass X-ray binaries (LMXBs) in equilibrium rotation, with the stellar magnetic field of order $B_s \\sim 10^{11}$ G. In contrast, van Paradijs et al. (1995) argue that these objects are isolated neutron stars accreting from a fossil disk, while Ghosh et al. (1997) suggest that AXPs are formed as the result of a Thorne-$\\dot{\\rm{Z}}$ytkov phase of a high mass X-ray binary with strong spherical accretion leading to the soft X-ray spectra with high foreground absorption. It is worth mentioning that accretion scenarios would be hard-pressed to explain the spin-down age of at least one member of Table 1, the $\\sim 2000$ yr of the pulsar in Kes~73 (see Table 1; Gotthelf \\& Vasisht 1997). First, it is difficult for accretion torques to spin-down a pulsar to 12-s in $\\sim 10^3$ yr from initial periods $P_i \\simlt 10^2$ ms unless, of course, the pulsar were born a very slow rotator, which is quite interesting in its own right.  Secondly, if the pulsar were rotating near its equilibrium period, as in the Ghosh and Lamb (1979) scenario, the spin-down time of ${P/ 2\\dot P} \\sim 3900$ yr is inconsistent with the luminosity implied accretion rate, $\\dot M \\simeq 10^{-11}$ M$_\\odot$ yr$^{-1}$ (assuming the pulsar has a standard dipolar field $\\simeq 10^{12}$ G); these usually lie in range $10^4 - 10^5$ yr.  In conclusion, further X-ray observations are required to secure the classification of the 7-s pulsar to the growing family of AXPs - by measuring the long term stability of the X-ray flux, secular trends in the pulse period including Doppler modulation, the lack of which will firm up the likelihood against an accreting binary hypothesis. Indeed, if AXPs are akin to SGRs then they could display sporadic hard X-ray transients, although such behavior is yet to be observed. Infrared observations to search for a possible counterpart to \\psr\\ could be carried out in spite the somewhat crude localization; we mention that past optical/IR searches for counterparts have been unsuccessful. Furthermore, radio observations to uncover an associated supernova remnant could be pursued. The high foreground absorption could easily cloak the soft X-ray emission of an aged, few$\\times 10^4$ yr-old, remnant. We find it remarkable that the period of \\psr\\ agrees so well with that of 1E~2259+586, although at present we believe this is sheer coincidence."
    },
    "9804/astro-ph9804213_arXiv.txt": {
        "abstract": "In efforts to demonstrate the linear Hubble law $ v = H r $ from galaxy observations, the underlying simplicity is often obscured by complexities arising from magnitude-limited data. In this paper we point out a simple but previously unremarked fact: that the shapes and orientations of structures in redshift space contain in themselves independent information about the cosmological redshift-distance relation.  The orientations of voids in the CfA slice support the Hubble law, giving a redshift-distance power index $ p = 0.83 \\pm 0.36 $ (void data from Slezak, de Lapparent, \\& Bijoui 1993) or $ p = 0.99 \\pm 0.38 $ (void data from Malik \\& Subramanian 1997). ",
        "introduction": "Hubble's (1929) observation that redshift increases linearly with distance for nearby galaxies has now been known for almost seventy years, and it is likely that its validity is not doubted by many. Yet, any attempt to demonstrate this simple law from galaxy observations soon descends into the complexities of magnitude-limited observations, the broad galaxy luminosity function, Malmquist bias, and, at larger redshifts, $K$-corrections, evolution, etc. This has permitted an often-ignored but persistent challenge to a linear redshift law from those preferring a quadratic relation, with both challengers and supporters citing data ranging from relatively nearby, bright optical galaxies (Soneira 1979, Segal 1980) through the {\\it IRAS} 1.2~Jy redshift catalog (Segal et al. 1993; Koranyi \\& Strauss 1997).  One complication in understanding the classical Hubble diagram has been inhomogeneity, appearing as a dependence of density on radius or redshift or in clustering and peculiar velocities. With the advent of more galaxy redshift catalogs covering larger amounts of the sky, we obtain ever clearer pictures of the universe. The CfA ``slice'' (de Lapparent, Geller, \\& Huchra 1986), in particular, first revealed dramatic structures, large voids separated by well-defined walls, extending to a significant fraction of the survey volume, perhaps even calling into question whether a survey to this depth represents a statistically fair sample of the universe. In other ways, however, this inhomogeneity itself can be useful. In this paper we use the shapes and orientations of structures to obtain information about the background cosmology in which they are embedded. In Section 2 below we discuss how shapes and orientations of structures in redshift space depend on the redshift-distance relation, and in Section 3 we apply these considerations to the CfA slice (de Lapparent, Geller, \\& Huchra 1986) and the Las Campanas redshift survey (Shectman et al 1996, LCRS). Section 4 contains a final discussion. ",
        "conclusions": "We have shown in this paper how the apparent shapes and orientations of objects in redshift space can be used to determine the redshift-distance relation. When redshift is not linear in distance, objects such as voids that are intrinsically round in space when viewed in redshift space are stretched along the line of sight; and an initially isotropic distribution of orientations becomes distorted in the radial direction, with measurable effect, as in \\eq{mu2d}. Application to voids in the CfA slice gives redshift power index $ p = 0.83 \\pm 0.36 $ from Slezak et al. (1993) and $ p = 0.99 \\pm 0.38 $ from Malik \\& Subramanian (1997), both a modest preference for $ p = 1 $ over $ p = 2 $. To obtain these results in redshift space with a true $ p = 2 $, voids in the CfA slice in space would have to be flattened and preferentially aligned transverse to the line of sight. The main complication to the simple interpretation is likely to come from peculiar velocities; an expansion velocity will introduce an apparent line-of-sight elongation. However, this effect is expected to be small: in the most extreme case, of an uncompensated, completely empty, isolated void in an $ \\Omega_0 = 1 $ universe, expansion would make $ p = 1 $ appear to be $ p = \\case43 $. In any case, unless voids are contracting, a condition difficult to make physical sense of, peculiar velocities will only increase, not decrease, the apparent value of $p$, and will not distort a quadratic redshift law to appear to prefer $ p = 1 $.  We do not expect void orientation statistics to replace classical methods for demonstrating the Hubble law. The most precise confirmation of the linearity of the Hubble law, $ p/5 = 0.2010 \\pm 0.0035 $, has been obtained recently from the Hubble diagram of type Ia supernovae (Riess, Press, \\& Kirshner 1996). The luminosity function as a function of redshift and the radial dependence of galaxy density in the {\\it IRAS} 1.2-Jy catalog (Koranyi \\& Strauss 1997) also support the linear Hubble law. However, our method offers an independent verification of the Hubble law in which magnitude-limited data, Malmquist bias, and evolution do not present problems. Further examination of the effect in real data, especially including in detail the peculiar velocities of the void walls, will undoubtedly introduce complications to the simple relations in equations (\\ref{mu3d}) and (\\ref{mu2d}). Whatever these complications, they will be different from those that enter arguments over the redshift-magnitude relation."
    },
    "9804/astro-ph9804163_arXiv.txt": {
        "abstract": "We present radio observations of the gamma-ray burster \\grb\\ made with the Very Large Array (VLA) and the Owens Valley Radio Observatory (OVRO) spanning a range of postburst timescales from one to 300 days. A search for a time-variable radio source was conducted covering an area which included a fading X-ray source and an optical transient, both of which are thought to be the long wavelength counterparts to the gamma-ray burst. At the position of the optical transient sensitive limits between 10 $\\mu$Jy and 1 mJy can be placed on the absence of a radio counterpart to \\grb\\ between 1.4 and 240 GHz. We apply a simple formulation of a fireball model which has been used with some success to reproduce the behavior of the optical and X-ray light curves.  Using this model we conclude that the radio non-detections are consistent with the peak flux density of the afterglow lying between 20-40 $\\mu$Jy and it requires that the optical flux peaked between 4 and 16 hours after the burst. ",
        "introduction": "The gamma-ray burst of 28 February 1997 was a turning point in our understanding of these enigmatic objects, resulting in the first-ever discovery of X-ray and optical counterparts to a burst.  Within the original 3-arcminute localization provided by the Wide Field Cameras (WFC) on board the BeppoSAX satellite, a previously unknown X-ray source 1SAX\\ts{J0501.7+1146} was detected by the Narrow Field Instruments (NFI). Eight hours after the burst the flux of 1SAX\\ts{J0501.7+1146} was 2.8$\\times{10}^{-12}$ erg cm$^{-2}$ s$^{-1}$ (2-10 keV), but three days later its flux had dropped by a factor of 20 (Costa et al. 1997). A comparison of V- and I-band images taken on 28 February and 8 March revealed an optical transient within a reduced error box, defined by the intersection of the WFC circle, the $\\pm$50\\arcsec\\ circle of the NFI, and the Interplanetary Network (IPN) annulus (van Paradijs et al. 1997, Hurley et al. 1997).  Predictions of long-lived afterglows from gamma-ray bursts at X-ray, optical and radio wavelengths have been made for some time (e.g. Paczy\\'nski \\& Rhoads 1993, M\\'esz\\'aros, Rees, \\& Papathanassiou 1994, Katz 1994, M\\'esz\\'aros \\& Rees 1997). In particular, a power-law decay in the observed long-wavelength flux with time is a generic consequence of a class of models known as ``fireballs''. The gamma-rays are thought to be produced when the relativistically expanding blast wave (aka fireball) is slowed down by the ambient gas or is self-shocked by its own ejecta. The fireball accelerates particles in a shock which then radiate via the synchrotron process (M\\'esz\\'aros et al. 1994, Waxman 1997a, 1997b, Sari, Piran, \\& Narayan 1998). As the fireball expands it cools, shifting the peak in the spectrum to lower energies and resulting in delayed emission at longer wavelengths.  The optical and X-ray decay from \\grb, as well as that from several other subsequent GRBs, is consistent at least to first order with one of the simpler formulations of these models (M\\'esz\\'aros \\& Rees 1997). Costa et al. (1997) fit a t$^\\delta$ decay to the X-ray data with $\\delta\\simeq-1.33\\pm{0.12}$, while global fits to the X-ray, optical and infrared data give $\\delta=-1.2$ (Wijers, Rees and M\\'esz\\'aros 1998) and $\\delta=-1.09\\pm0.23$ (Reichart 1997). More recent optical fits, aided by a long time-baseline, yield $\\delta=-1.12\\pm0.08$ (Garcia et al. 1997), $\\delta=-1.21\\pm0.02$ (Masetti et al. (1997), and $\\delta=-1.10\\pm0.04$ (Galama et al. 1998). The character of this decay agrees well with a fireball produced by a one-time impulsive injection of energy in which only the forward blast wave efficiently accelerates particles (i.e. the adiabatic piston model). The adiabatic model predicts a simple relation between the slope of the temporal decay $\\delta$ and the flux spectrum $\\beta$=${2\\over3}\\delta$. Observationally the spectrum is not well determined, with $\\beta=-0.7\\pm0.1$ optically (van Paradijs et al. 1997) and X-ray values of $\\beta\\simeq-0.9$ with large scatter (Frontera et al. 1998).  The slope of the particle spectrum $p$ (where $\\beta=(p-1)/2)$ inferred from these values of temporal decay is $p\\sim-2.6$, not an unreasonable value for a relativistic shock (Blandford and Eichler 1987).  Given the early success of the fireball model in predicting the gross properties of the optical and X-ray behavior of \\grb, it is reasonable to look for delayed radio emission from this burst. This is particularly relevant in the light of the discovery of the radio afterglow from GRB\\thinspace{970508} (Frail et al. 1997a) which exhibited temporal and spectral behavior consistent with this model (Waxman, Kulkarni \\& Frail 1997).  The properties of the fireball for \\grb\\ are well constrained by the optical and X-ray data.  Thus the presence or absence of radio emission at late times has a bearing on the validity of this model. With this in mind, we began a radio search centered on the optical transient detected by van Paradijs et al. (1997). This {\\it Letter} is a summary of our monitoring program for the first 300 days. ",
        "conclusions": "The detection of an optical transient, thought to be the afterglow from \\grb, has made it possible to perform a search for time-variable radio emission at the optical position. Delayed emission at longer wavelengths (X-ray, optical and radio) is a generic prediction of all fireball models. VLA and OVRO observations have been made that span a range of postburst timescales from one to 300 days putting limits on the absence of a radio counterpart to \\grb\\ between 10 $\\mu$Jy and 1 mJy.  Applying a simple version of the fireball model which has been used successfully to fit the temporal behavior of the decaying X-ray and optical emission from \\grb\\ suggests that the radio afterglow has yet to peak. If this is correct then continued deep imaging in the coming months offers the possibility for the detection of a weak but increasing radio signal. Detecting this emission would be a powerful verification of the fireball model as described by M\\'esz\\'aros \\& Rees (1997).  Indeed, as the recent detection of GRB\\thinspace{970508} has taught us, radio afterglows yield unique GRB diagnostics that are not obtainable by any other means (Frail et al. 1997b).  Unlike optical or X-ray wavelengths one is presented with the possibility of following the {\\it full} evolution of the fireball emission through all its different stages; first while it is optically thick, then as it slowly rises to a peak flux density and thereafter decays, making a transition from an ultra-relativistic to sub-relativistic shock. Furthermore, both the scintillation of the radio source (Goodman 1997) and its flux density, when it is synchrotron self-absorbed (Katz 1994), allow a determination of the size and expansion of the fireball.  If the radio afterglow is going to be detected from \\grb\\ and others like it, then continued long-term monitoring is going to be required at the microJansky level."
    },
    "9804/astro-ph9804040_arXiv.txt": {
        "abstract": " ",
        "introduction": "Few properties of astronomical masers are determined directly by observations, most are inferred indirectly.  Foremost among the latter is the maser saturation stage. Saturation has a significant impact on maser growth, so determining whether a maser is saturated ($J > J_s$, where $J$ is the angle-averaged intensity and $J_s$ is the saturation intensity) or not is usually a precondition for analysis of the observations.  Unfortunately, this crucial issue is not convincingly settled.  Strong masers are generally believed to be saturated, but in most cases the evidence is less than compelling as it relies primarily on plausibility arguments rather than quantitative tests (see e.g.\\ \\[Book], \\S 8.6).  This unsatisfactory situation reflects a fundamental difficulty --- neither $J$ nor $J_s$ is directly measurable.  The saturation parameter $J_s$ is a theoretical quantity, determined only within the context of a given pumping scheme.  And because maser radiation is highly beamed, $J = I\\Omega/4\\pi$ so this quantity, too, cannot be directly measured; the intensity $I$ is measurable when the maser is resolved, but the beaming angle $\\Omega$ is unobservable.  Similarly, the amplification optical depth has never been directly determined for any maser that amplifies its own radiation.  Recent VLA observations of OH 1720 MHz masers near the Galactic center by \\[YZ96] open up new possibilities for direct determination of some maser properties. Significant circular polarization (upward of 20\\%) is detected in various spectral features, and the right- and left-hand components coincide on the sky, as expected from the Zeeman effect. Furthermore, the spectral shape of the Stokes parameter $V$ follows an antisymmetric S-curve with sharp reversal at line center, the typical profile for Zeeman shift \\DnuB\\ much smaller than the Doppler linewidth \\DnuD.  Similar results were previously reported for H$_2$O masers in star-forming regions by \\[FbG] and for OH 1612 MHz masers in OH/IR stars by \\[ZFix], but the polarization was lower and the quality of the data not nearly as high.  The general maser polarization solution was recently derived for arbitrary values of \\eq{ \\xb = {\\DnuB \\over \\DnuD} } (Elitzur 1996, hereafter \\[E96]) and the solution properties at $\\xb \\ll 1$ closely match the observed circular polarization. Here I show that a comparative analysis of the spectral profiles of $I$ and $V$, two measurable independent maser intensities, offers direct determination of various maser properties, in particular the saturation stage.  The analysis is readily performed with the aid of the ratio profile \\eq{\\label{R0} \\R(\\nu) = {V(\\nu) \\over I'(\\nu)} = {v(\\nu)\\over I'(\\nu)/I(\\nu)} } where the prime denotes derivative with respect to frequency and $v = V/I$ is the fractional circular polarization.  When $\\xb \\ll 1$, spectral analysis of \\R\\ offers intrinsic sensitivity of order \\xb\\ and has long been an important tool in studies of the Zeeman effect of thermal radiation (see e.g.\\ \\[Trol]). In that case \\R\\ is constant across the spectral line and its magnitude determines the magnetic field along the line of sight.  This constancy of \\R\\ follows from some simple, general symmetry arguments as shown by \\[Cru] (see also \\S3 below). However, maser exponential amplification during unsaturated growth destroys both the underlying symmetry and the constancy of \\R, the saturation process restores both. The key to the different behavior, and \\R-profiles, in the two regimes is the narrowing of the maser line during unsaturated amplification and its rebroadening during saturation.  The important differences between thermal and maser polarization are discussed in detail below.  For completeness, some basic elements of the polarization theory developed in \\[E96] are reproduced in \\S2.  The \\R\\ profile is discussed in \\S3 for thermal radiation and in \\S4 for maser radiation when $\\xb \\ll 1$. In \\S5, circular polarization for fully resolved Zeeman patterns, $\\xb > 1$, is discussed.  The implications for observations are discussed in detail in \\S6. ",
        "conclusions": ""
    },
    "9804/astro-ph9804089_arXiv.txt": {
        "abstract": "Stellar dynamics is almost unreasonably well suited for an implementation in terms of special-purpose hardware.  Unlike the case of molecular dynamics, stellar dynamics deals exclusively with a long-range force, gravity, which leads to a computational cost scaling as the square of the number of stars involved.  While special tricks can lead to a reduction of this cost from $\\sim N^2$ to $\\sim N\\log N$ in the case of very large particle numbers, such tricks are not suitable for all areas within stellar dynamics.  When a stellar system is close to equilibrium, and has a very high density, it still pays to compute all interactions on a star by star basis, even for $N=10^5$. Any $cN\\log N$ approach would either gloss over the subtle net effects of near-canceling interactions, driving the evolution of such a system, or would carry a prohibitively large coefficient $c$.  This paper presents a brief introduction to the stellar dynamics of dense stellar systems, aimed at researchers using special purpose computers in other branches of physics. ",
        "introduction": "Stellar dynamics is the branch of astrophysics that studies the structure and evolution of collections of stars, from small groups to larger star clusters to entire galaxies and clusters of galaxies.  The interactions between the individual stars can be modeled to a high degree of accuracy as Newtonian gravitational interactions between point masses.  Only under extremely high densities, such as occurs in the nuclei of galaxies and the centers of the densest star clusters, do stars have a reasonable chance to undergo a physical collision during their life time.  In contrast, a typical star, such as our own Sun, has a probability of only 1 in $10^{8}$ to undergo a collision with a neighboring star, during the remaining $5\\times10^9$ years of its life.  The most spectacular example of a dense stellar system within our own galaxy is the agglomeration of roughly a million stars within the inner parsec from the center (a parsec is a unit of length, equal to a few light years, and corresponds to a typical distance between stars in the solar neighborhood).  These stars describe orbits around the black hole that resides in the very center of our galaxy.  The black hole itself has a mass that is a few million times larger than the mass of the Sun.  The density of stars around the black hole is a million times larger than the stellar density in a typical part of the galaxy, such as where we reside.  A detailed stellar dynamical modeling of the center of our galaxy is still difficult, partly because the observations of this heavily obscured area have only recently become accurate enough to tell us the physical characteristics of the system, partly because of the interference of other physical effects, such as the presence of gas clouds and ongoing star formation.  Before tackling the stellar dynamics of the nucleus of our galaxy, it is therefore prudent to start our attempts with a simpler system, such as is provided by the core of a dense globular cluster.  While most of the stars in and around our galaxy are spread out throughout the disk, and to a lesser extent through the halo, there are more than one hundred isolated star clusters circling the galaxy, each containing of order $10^6$ stars.  In a dozen or so of those globular clusters, as they are called because of their appearance, the central densities rival that of the density in the nucleus of our galaxy.  However, the absence of a large black hole, as well as gas clouds and concomitant star formation, makes it far easier to study and model globular cluster cores in detail.  In addition, recent observations, notably with the Hubble Space Telescope ({\\it cf.} \\cite{Guh96}), have resolved those cores into individual stars, something that has not been possible with ground-based observations.  This paper sketches some of the progress made in the study of globular cluster cores, emphasizing the role of special-purpose computers. ",
        "conclusions": ""
    },
    "9804/astro-ph9804276_arXiv.txt": {
        "abstract": "Modelling gravity is a fundamental problem that must be tackled in $N$-body simulations of stellar systems, and satisfactory solutions require a deep understanding of the dynamical effects of softening. In a previous paper (Romeo 1997), we have devised a method for exploring such effects, and we have focused on two applications that reveal the dynamical differences between the most representative types of softened gravity. In the present paper we show that our method can be applied in another, more fruitful, way: for developing new ideas about softening. Indeed, it opens a {\\it direct\\/} route to the discovery of optimal types of softened gravity for given dynamical requirements, and thus to the accomplishment of a physically consistent modelling of disc galaxies, even in the presence of a cold interstellar gaseous component and in situations that demand anisotropic resolution. ",
        "introduction": "$N$-body simulations of disc galaxies rely on the use of softening. This artifice removes the short-range singularity of the gravitational interaction, which is dynamically unimportant and computationally troublesome, whereas it leaves the long-range behaviour of gravity unchanged. But softening is also a critical factor in simulations. It controls their quality and can affect their result on scales much larger than the softening length. Its dynamical effects are further exacerbated in the presence of a cold interstellar gaseous component and in situations that demand anisotropic resolution. Thus softening poses a dynamical problem of special concern, which should be probed carefully and in detail (e.g., Hernquist \\& Barnes 1990; Pfenniger \\& Friedli 1993; Romeo 1994, hereafter Paper I; Romeo 1997, hereafter Paper II% \\footnote{Sections and equations of that paper are denoted by the prefix II.}; and references therein).  In Paper I, we have investigated how faithful simulations are. In particular, we have concluded that the standard way of introducing softening in the presence of stars and cold interstellar gas is definitely unsatisfactory in several regimes of astrophysical interest. It is so because important small-scale instabilities of the gaseous component, e.g.\\ those peculiar to star-formation processes, are suppressed just as unphysical noise of the stellar component. Faithfulness requires an appropriate introduction of two softening lengths, one for each component, and also a rigorous specification of the star-gas gravitational interaction.  In Paper II, we have devised a method for exploring the dynamical effects of softening. As a major result, we have shown how to choose the softening length for optimizing the faithfulness of simulations to the Newtonian dynamics. Then we have focused on two applications that reveal the dynamical differences between the most representative types of softened gravity. In particular, we have concluded that it is desirable to improve the current way of introducing anisotropic softening. We need a clearer decoupling of the resolution parallel and perpendicular to the plane, and also more natural planar and vertical softening lengths.  In the present paper, which completes our planned research work about softening, we propose an {\\it innovative\\/} solution to the problem. The understanding of galactic and extragalactic astrophysics is at a crucial stage. Unsolved problems are viewed in new perspectives, which promise major revisions of knowledge (see, e.g., Blitz \\& Teuben 1996; Block \\& Greenberg 1996). Recent investigations suggest, for instance, a more enigmatic interplay between stellar disc and bulge/halo (e.g., Lequeux et al.\\ 1995), a clearer relation between cold gas and dark matter in spiral galaxies (e.g., Pfenniger et al.\\ 1994; Pfenniger \\& Combes 1994; Combes \\& Pfenniger 1997), and a closer connection between the fractal structures of the interstellar medium and of the universe (e.g., de Vega et al.\\ 1996, 1998). The implications are clear: modelling gravity in $N$-body simulations of disc galaxies should offer a flexible interface with such a progress. Our solution is to optimize the fidelity of simulations to given dynamical requirements. How do we apply this idea in practice? \\begin{enumerate} \\item We impose the requirements in the wavenumber space since this is the natural dynamical domain of gravity, as Pfenniger \\& Friedli (1993) have previously emphasized. \\item We identify the softening length with the characteristic dynamical scale length. \\item Then we invert part of the method of Paper II, and the result is the optimal type of softened gravity that satisfies those dynamical requirements. \\end{enumerate} Our application covers both 2-D and 3-D modelling. The basic cases are extended to more complex situations through recipes for implementing star-gas and anisotropic softening, which have already been motivated (cf.\\ discussions of Papers I and II). Last but not least, each description is complemented by an example that leaves room for creativity.  The present paper is organized as follows. The application is shown in Sects.\\ 2 and 3 (see also Appendix A), and proceeds as in the previous discussion. Comments on related works concerning softening are made in Sect.\\ 4. The conclusions and perspectives are drawn in Sect.\\ 5, where we present our three papers about softening in a more unified view and emphasize their potentially strong impact on galactic dynamics.   \\begin{figure*} \\vbox{\\vspace{.1cm} \\hbox{\\hspace{-.25cm} \\psfig{figure=romeof1.ps,width=18.9cm,height=13.275cm,angle=-90}} \\vspace{-2.9cm}} \\hfill\\parbox[b]{5.7cm}{\\caption[]{Examples of 2-D modelling: {\\bf a} one-com\\-pon\\-ent case (cf.\\ Sect.\\ 2.1), {\\bf b} two-com\\-pon\\-ent case (cf.\\ Sect.\\ 2.2). The abbreviations N, T and P mean Newtonian gravity, thickness and Plummer softening, respectively}} \\end{figure*}   \\begin{figure*} \\vbox{\\vspace{.1cm} \\hbox{\\hspace{-.25cm} \\psfig{figure=romeof2.ps,width=18.9cm,height=13.275cm,angle=-90}} \\vspace{-2.9cm}} \\hfill\\parbox[b]{5.7cm}{\\caption[]{Examples of 3-D modelling: {\\bf a} iso\\-tropic case (cf.\\ Sect.\\ 3.1), {\\bf b} an\\-iso\\-tropic case (cf.\\ Sect.\\ 3.2). The abbreviations N, T and P mean Newtonian gravity, thickness and Plummer softening, respectively}} \\end{figure*} ",
        "conclusions": "The importance of computer simulations in astrophysics is analogous to that of experiments in other branches of physics. They also serve as a welcome bridge between theories, often restricted to idealized situations, and observations, revealing instead the complexity of nature. Major present objectives are to construct physically consistent $N$-body models of disc galaxies and to simulate their dynamical evolution, especially in regimes of spiral structure in which a fruitful comparison between theories and simulations can be made (e.g., Pfenniger \\& Friedli 1991; Junqueira \\& Combes 1996; Zhang 1996; Bottema \\& Gerritsen 1997; Fuchs \\& von Linden 1998; von Linden et al.\\ 1998; Zhang 1998a, b). The construction of such models is indeed a difficult task which has not yet been fully accomplished, and which should eventually provide clues of vital importance to a number of open questions posed by both theories and observations.  Our involvement has been threefold. In Paper I, we have recognized a fundamental problem posed by this research programme (for a concrete use of that analysis and for interesting remarks see, e.g., Junqueira \\& Combes 1996). In Paper II, we have devised a method for solving this problem. In the present paper, we apply this method and solve the problem, thus laying the foundations of such a plan. The {\\it major result\\/} is that gravity can be modelled so as to optimize the fidelity of simulations, and the procedure is practicable. The following conclusions point up the whys and wherefores: \\begin{enumerate} \\item Optimization is performed with respect to arbitrary dynamical requirements and, in specific examples, with respect to the Newtonian dynamics. This enriches the modelling with an {\\it unprecedented\\/} degree of freedom, which has clear epistemological motivations (cf.\\ Sect.\\ 1, discussion of the present paper). \\item Optimization is performed in the wavenumber space. This is the {\\it appropriate\\/} domain for imposing dynamical requirements on the modelling. \\item Optimization concerns {\\it both\\/} the softening length {\\it and\\/} the type of softened gravity. \\item Softening is conceived as a {\\it double\\/} artifice. The softened gravity and finite-sized particle conceptions are equivalent in the basic cases. Concerning more complex situations, the latter is particularly useful for implementing star-gas softening, whereas the former is particularly useful for implementing anisotropic softening. Thus both conceptions contribute towards the accomplishment of a physically consistent modelling. \\end{enumerate} Our application is ready for a concrete use. An attractive idea is to employ a particle-particle code together with MD-GRAPE, a highly parallelized special-purpose computer for many-body simulations with an arbitrary central force (Fukushige et al.\\ 1996). We can also employ a classical particle-mesh code. Then the dynamical effects of the grid are known and factorize as those of softening (e.g., Bouchet et al.\\ 1985; Efstathiou et al.\\ 1985; for a review see, e.g., Hockney \\& Eastwood 1988). So essentially the application proceeds as in the present paper, but it may be useful to act directly on the wavenumber space (e.g., Tormen \\& Bertschinger 1996). A more complex problem concerns tree codes, which have hierarchical structure and adaptive resolution over multiple scales (e.g., Hernquist 1987; for a review see, e.g., Pfalzner \\& Gibbon 1996). The solution to that problem would need a more advanced analysis (cf.\\ following discussion). Welcome suggestions about the choice of the code can come from cosmological simulations (e.g., Splinter et al.\\ 1998).  Finally, what about the future? Our approach is connected with the technique of filtering in spectral domain used in the context of digital image processing. This is a rapidly evolving field with growing applications in science and engineering, which can promote further substantial advances in $N$-body modelling of disc galaxies. For instance, wavelets are ideal for resolving multi-scale problems in space and/or time, such as those concerning turbulence, bifurcations, fractals and many others (see, e.g., Kaiser 1994; Holschneider 1995; Bowman \\& Newell 1998; for an alternative analysis tool see, e.g., Stutzki et al.\\ 1998). Speculating further, wavelets might be used for speeding up simulations through fast solution of linear systems (cf.\\ Press et al.\\ 1992, pp.\\ 597--599 and 782).  These are the merits of our contribution. We hope that the trilogy (Papers I--III) and further reflections (Romeo 1998) will encourage $N$-body experimenters to model gravity so as to optimize the fidelity of their simulations, and that the result will be a stronger interdisciplinary connection with theories and observations."
    },
    "9804/astro-ph9804106_arXiv.txt": {
        "abstract": "We propose that the majority of quasars at redshift $z\\sim 1 - 5$ formed in the environment of new born collapsed halos with 1-D velocity dispersion $\\sigma_v^{1d} \\sim 400 \\kms$. The harboring coefficient $f$ of quasars per halo and the lifetime of quasars depend only on local process, not modulated by the density inhomogeneities on scales larger than the size of the halos. Thus, the bias of quasars on scale larger than the size of these halos is mainly determined by the parameter $\\sigma_v$ used for quasar environment identification. With this model, the popular structure formation models, like SCDM and LCDM, can be fairly well reconciled with the data of quasars, including a. observed feature of the environment for quasars; b. redshift evolution of quasar abundance; c. the two-point correlation functions of quasars. This bias model predicts that the correlation function of quasars doesn't significantly evolve, or only slightly increases with redshift. ",
        "introduction": "\\bigskip  Mass distribution at high redshifts is being a hot subject of the large scale structure study. Data of various objects at moderate and high redshifts, in particular, clusters of galaxies, are becoming available for probing the formation and evolution of structures and for discriminating among popular dark matter models (e.g. Jing \\& Fang 1994; Eke, Cole \\& Frenk 1996; Bahcall, Fan \\& Cen 1997; Kitayama \\& Suto 1997). Considering that quasars are the most distant among various luminous objects, they have also been applied in this approach (e.g. Bi \\& Fang 1997).  However, as a mass tracer of the cosmic matter field, quasars are still playing a role different from clusters of galaxies. The problem stems from so-called ``bias''. Clusters are a biased tracer of the mass distribution. The correlation amplitude of clusters is believed to be much higher than that of the underlying matter and increases strongly with the cluster richness (Bahcall \\& Soneira 1983). This bias is plausibly explained by the mechanism that the observed clusters are identified as massive collapsed halos of the density field (Kaiser 1984).  That is, the bias of clusters is completely determined by the gravitational parameters, like mass $M$ and virial radius $r_{vir}$ used to identify the halos.  With this approach, a detailed confrontation can be made between theories and the observations of clusters.  Quasars may also be biased tracer of the mass distribution. Recent observations indicate that the correlation amplitude of quasars may also be different from the underlying dark matter (Mo \\& Fang 1993; Komberg, Kravtsov \\& Lukash 1994; Croom \\& Shanks 1996; Franca, Andreani \\& Cristiani, 1997).  However, so far no detailed model is available for the bias of quasars, though their high clustering strength and environment imply that quasars are hosted by massive halos (see below). Because of the lack of such a model, one cannot confront the data of quasars with theoretical models in the way as for clusters. For instance, the abundance of quasars can only be used as an upper or lower limit to certain massive halos; no detailed comparison between the number densities of quasars and of halos is allowed. Obviously, it is very important to understand what kind of mass halos are associated with the majority of quasars. Such a knowledge will not only enable the observational data of quasars to be powerful tests for theoretical models of galaxy formation but also tell that what type of local environments is responsible for intriguing the nuclear activities of quasars.  Like clusters and groups of galaxies, it is generally believed that quasars should be associated with certain type of collapsed dark matter halos. Yet, different from identification of clusters, the environment suitable for forming quasars is not merely determined by gravitational parameters, as the hydro processes are also involved. Therefore, the identification of quasar-harboring halos should be given by both gravitational and hydro parameters. In other words, not all halos with certain $M$ and/or velocity dispersion $\\sigma_v$ harbor quasars, because certain hydro conditions must be satisfied. However, considering the hydro processes are local, it is reasonable to assume that the hydro conditions may not be modulated by the density inhomogeneities on scales much larger than the size of the halo $l$. In this case, the probability for a halo to have a quasar should be the same for all halos of the same kind, without depending on structures larger than $l$. Thus, the relative fractions of quasars with respect to the certain collapsed halos should be the same for all volumes larger than $l^3$. Consequently, when averaged on scale larger than $l$, the distribution of quasars $n_{qso}({\\bf r},z)$ at redshift $z$ should be proportional to that of the considered halos, $n_{halo}({\\bf r},z)$. Thus, all effects of the hydro processes can be absorbed into a ``normalization factor\" $A$, i.e. $n_{qso}({\\bf r},z)=An_{halo}({\\bf r},z)$, and $A$ is less dependent on $z$ than $n_{halo}({\\bf r},z)$. The bias of quasar distribution with respect to the mass distribution is then dominated by the bias of the selected halos with respect to the mass. Based on this analysis, quasar bias, at least on large scales, may also be only gravitational, depending on the gravitational parameters used for selecting the quasar-suitable halos.  Accordingly, a possible model for quasar bias should at least satisfy the three conditions. 1. Gravitational environment given by the identified halos is consistent with the observed environment around quasars; 2. The abundance of quasars, $n_{qso}(z)$, at redshift $z$ is proportional to the number density of the identified halos, $n_{halo}(z)$ in a redshift-independent way, i.e. $n_{qso}(z)=An_{halo}$ where $A$ is a {\\it z-independent} constant, 3. The amplitude and $z$-evolution of the halo-halo correlation function are consistent with the observed correlation function of quasars. In this letter, we will show within the framework of the CDM cosmogonic theories that such a bias model can indeed be settled following the above-mentioned points.  The details of the points 1, 2 and 3 will be discussed in the \\S 2, 3 and 4, respectively. ",
        "conclusions": "We showed that velocity-dispersion-selected halos are a possible mechanism for the bias of quasars. The majority of quasars at redshift $z\\sim 1 - 5$ formed in the environment of new born collapsed halos with 1-D velocity dispersion $\\sigma^{1d}_v \\sim 400 \\kms$.  Both the harboring coefficient $f$ per halo and the lifetime of quasars are $z$-independent. With this bias model, the popular structure formation models, like SCDM and LCDM, can be fairly well reconciled with data of the abundance and correlations of quasars at $z \\geq 0.5$.  It is interesting to point out that the velocity dispersion identified halos generally don't have the same mass. Eq.(2) shows that for a given $\\sigma^{1d}_v$, the redshifts the higher, the mass of the halos the smaller. This result has already been recognized in an earlier study, which shows that in order to fit with quasar abundance at high redshift, the mass of the halos has to be smaller than at the lower redshift (Bi \\& Fang 1997).  With this model, one can predict that 1. The environment for quasars at redshifts from $z \\sim 1$ to 5 should be characterized by a velocity dispersion, $\\sigma^{1d}_v \\sim 400 \\kms$; 2. The amplitudes of quasar two-point correlation function at high redshifts don't significantly evolve with redshifts. In the paper, only the models of the SCDM and LCDM are considered. We can expected that with better data of quasars becoming available, the bias model of quasars will play more important role for discriminating among models of structure formations."
    },
    "9804/astro-ph9804330_arXiv.txt": {
        "abstract": "s{We review present understanding of Galactic free--free emission and its possible importance to CMB fluctuation measurements. Current results, from both ``direct'' observations in the microwave band and from H$\\alpha$ studies, suggest that this foreground does not represent a serious obstacle to mapping the CMB; however, this is based on limited information and we emphasize the need for more exhaustive studies. We also present some preliminary results based on our recent H$% \\alpha$ observations near the South Pole CMB data sets. The fluctuation amplitude seen in H$\\alpha$ indicates that the detected CMB fluctuations are not significantly contaminated by free--free emission, at least if the diffuse gas is at a temperature of $T\\sim 10^4$ K.} ",
        "introduction": " ",
        "conclusions": "To summarize the current status of our understanding of Galactic free--free emission, we would say that although there is {\\em no indication} of fluctuations large enough to pose serious difficulties for CMB observations, the observational constraints remain weak. A critical interpretation of the results in the tables would be that the limit on large scale free--free fluctuations is the same order as the CMB amplitude on these same scales (at $~40$ GHz). On smaller scales, observations in $H\\alpha$ have not turned up any signs of large amplitude variations, but those based on high resolution spectrographs are few and cover only a small percentage of the sky. There does appear to be a dust/free--free correlation on all angular scales, but there is room, and perhaps tentative indications of, an equally important non--correlated component (question \\#1 posed in the introduction). And then there is the puzzeling result from Leitch et al. (1997), perhaps pointing to a hot phase of the ISM which could, due to lack of sensitivity, escape many of the present $H\\alpha$ limits (question \\#2 posed in the introduction).   Obviously, CMB observations at higher frequencies, where much of the effort is now being concentrated, will suffer less from any possible free--free contamination, and the many of the next generation CMB experiments have a wide spectral coverage to aid the removal of foregrounds.  Even given the above critical viewpoint, it would be a surprise to discover free--free emission presenting an important difficulty for all planned CMB experiments, at least in terms of measuring the variance, or power spectrum, from the early Universe.  Foregrounds will, however, be much more important for higher order statistics looking for non--gaussian signatures.  In such cases, the non--gaussian foregrounds will have to be removed to high precision.   A sensitive, high spectral resolution H$\\alpha$ survey of the entire sky would be of great value in the context of the above considerations.  We have also reported some results from our recent H$\\alpha$ observations taken with a Fabry--Perot system at La Silla. The data cover two bands along which the South Pole data sets show significant fluctuations in CMB brightness. The observed H$\\alpha$ fluctuations indicate that the CMB results in this region are not significantly contaminated by free--free emission (assuming a gas with $T_4=1$)."
    },
    "9804/astro-ph9804218_arXiv.txt": {
        "abstract": "We present new spectroscopic and photometric time series observations of the $\\delta$ Scuti star FG~Vir.  We detect the oscillations via changes in the equivalent widths of hydrogen and metal absorption lines.  {}From the ratios between spectroscopic and photometric amplitudes, we assign $\\ell$ values to the eight strongest oscillation modes.  In particular, we identify two radial modes ($\\ell =0$) and find that the main pulsation mode (147~$\\mu$Hz) has $\\ell =1$.  One of the radial modes (at 140$\\mu$Hz) is the fundamental, implying that two modes with lower frequencies are {\\it g}-modes.  For the radial modes, we compare frequencies with those calculated from a scaled $\\delta$~Scuti star model and derive a density $0.1645\\pm 0.0005\\,\\rho_{\\odot}$.  We then obtain a distance of $84\\pm 3$\\,pc, in excellent agreement with the Hipparcos value.  Finally, we suggest that a 3.5-day variability in all observables (equivalent widths and intensity) is caused by stellar rotation. ",
        "introduction": "\\label{intro}  Oscillations in multi-periodic variables such as $\\delta$ Scuti, roAp and $\\beta$ Cephei stars have been observed extensively during the past 20 years.  But even with high-quality data, it is still extremely difficult to identify which modes are being detected.  Kjeldsen et al.\\ (\\cite{kbvf95}) used a new technique to detect solar-like oscillations in the bright G sub-giant $\\eta$ Boo through their effect on the equivalent widths of the Balmer lines.  A subsequent discussion by Bedding et al.\\ (\\cite{bkrb96}) of the sensitivity of different observables to modes with different degree $\\ell$ suggested that one can determine the $\\ell$-value of a given mode by combining measurements of absorption-line equivalent widths with simultaneous photometric observations.  To test this idea, we chose the bright and well-studied $\\delta$~Scuti star FG~Vir.  FG Vir (HD 106384; $V=6.57$) is a multi-periodic $\\delta$~Scuti star.  It has a main pulsation period close to 1.9 hours and shows a fairly complex oscillation spectrum.  This star has been studied extensively during the last few years, resulting in the detection of at least 24 well-determined frequencies between 100 and 400~$\\mu$Hz, with amplitudes from 0.8 to 22 milli-magnitudes (mmag; Breger et al.\\ \\cite{bhn95}, \\cite{br98}).  Because of its slow rotation (which reduces the complicating effects of rotational splitting) and the large number of detected frequencies (some of which are probably {\\it g} modes), FG Vir is one of the most promising candidates for performing asteroseismology on a $\\delta$ Scuti variable.  Observations and models of this star have been presented by Dawson et al.\\ (\\cite{dbl95}), Breger et al.\\ (\\cite{bhn95}) and Guzik \\& Bradley (\\cite{gb95}).  By choosing a star like FG Vir we have the advantage of knowing the frequencies in advance. We are therefore able to determine the oscillation amplitudes and phases with high precision.  The aim of this paper is to identify the $\\ell$ values of the observed modes and to compare the oscillation frequencies with a pulsation model. A preliminary analysis of the observations presented in this paper was given by Viskum~et~al.~(\\cite{vdk97}), while results on radial velocity measurements were given by Viskum et al.\\ (\\cite{vb97}). ",
        "conclusions": "We have investigated a new technique to measure the oscillations in $\\delta$~Scuti stars via changes in the equivalent widths of absorption lines.  An important advantage of this new technique is that only medium-dispersion spectra are needed, which makes the method suitable for small and medium-sized telescopes and for multi-site campaigns.  Our main results are summarized below. \\begin{itemize}  \\item Our detection of oscillations in FG~Vir from equivalent-width measurements of \\Ha, \\Hb\\ and Fe\\,{\\sc i} lines is an important confirmation of the method developed by Kjeldsen et al.\\ (\\cite{kbvf95}) to search for solar-like oscillations.  \\item {}From the ratios between oscillation amplitudes measured in EW and the four Str\\\"omgren filters ({\\it uvby}), we have identified $\\ell$-values for eight modes in FG~Vir.  \\item We suggest that the two lowest-frequency modes are {\\it g}-modes, while the strongest mode (147.2$\\mu$Hz) is a dipole mode.  \\item {}By comparing the frequencies of radial modes with model calculations, we obtained a precise density and derived a distance that is in excellent agreement with the Hipparcos value.  \\item We detected a long-period variation in the time series with a period of 3.5 days, which we propose is caused by rotation of the star.  \\end{itemize}"
    },
    "9804/astro-ph9804112_arXiv.txt": {
        "abstract": "We discuss a technique for mapping the synchrotron turnover frequency distribution using nearly simultaneous, multi--frequency VLBI observations. The limitations of the technique arising from limited spatial sampl\\-ing and frequency coverage are investigated. The errors caused by uneven spatial sampl\\-ing of typical multi--frequency VLBA datasets are estimated through numerical simulations, and are shown to be of the order of 10\\%, for pixels with the deconvolution ${\\rm SNR} \\sim 7$. The fitted spectral parameters are corrected for the errors due to limited frequency coverage of VLBI data. First results from mapping the turnover frequency distribution in \\object{3C\\,345} are presented.  \\keywords {methods: data analysis -- methods: observational -- quasars: individual: \\object{3C\\,345}} ",
        "introduction": "}  Information obtained with Very Long Baseline Interferometry (VLBI) about radio spectra of parsec--scale jets and their evolution can be crucial for distinguishing between various jet models. However, there are several aspects of VLBI which impede spectral studies of parsec--scale regions.  The reliability of spectral information extracted from VLBI data depends on many factors including sampl\\-ing functions at different frequencies, alignment of the images, calibration and self--calibration errors, a narrow range of observing frequencies, and source variability.  The influences of all these factors must be understood and, if possible, corrected for, in order to reconstruct the spectral properties of parsec--scale jets consistently.  Radio emission from the parsec--scale jets is commonly described by the synchrotron radiation from a relativistic plasma (e.g. Pacholczyk 1970). The corresponding spectral shape, $S(\\nu) \\propto \\nu^{\\alpha}$, is characterized by the location of spectral maximum ($S_{\\rm m}$, $\\nu_{\\rm m}$) also called the turnover point, and by the two spectral indices, $\\alpha_{\\rm thick}$ (for frequencies $\\nu \\ll \\nu_{\\rm m}$) and $\\alpha_{\\rm thin}$ (for $\\nu \\gg \\nu_{\\rm m}$).  In many kiloparsec--scale objects, spectral index distributions have been mapped, using observations made with scaled arrays. In such observations, the antenna configurations are selected at each frequency in a specific way such that the spatial sampl\\-ings of the resulting interferometric measurements are identical at all frequencies used for the observations.  It is virtually impossible to use the scaled array technique for VLBI observations of parsec--scale jets made at different frequencies. The uneven spatial sampl\\-ings of VLBI data at different frequencies result in differences of the corresponding synthesized beams, and can ultimately lead to confusion and spurious features appearing in spectral index maps.  In spectral index maps, the only available kind of information is the spectral slope between the two frequencies. While sufficient for many purposes, this information can be misleading in the situation when the frequency of thespectral maximum lies between the frequencies used for spectral index mapping.  In the ranges of frequencies between 1.4 and 43\\,GHz, frequently used for VLBI observations, such a situation can be quite common. Using observations at three or more frequencies, it is possible to estimate the shape of the synchrotron spectrum, and derive the turnover frequency (frequency of spectral maximum). Information about the turnover frequency can help to avoid the confusion which is likely to occur in spectral index maps.  The turnover frequency is sensitive to changes of physical conditions in the jet such as velocity, particle density, and magnetic field strength. This makes it an excellent tool for probing the physics of the jet in more detail than is allowed by analysis of the flux and spectral index properties of the jet.  In this paper, we present a technique suitable for determining the turnover frequency distribution from multi--frequency VLBI data, and investigate its limitations and ranges of applicability. We discuss the advantages of using the Very Long Baseline Array\\footnote{The Very Long Baseline Array is operated by the National Radio Astronomy Observatory (NRAO)} (VLBA) for spectral imaging.  A general approach to imaging of VLBA data from nearly simultaneous, snapshot--type observations at different frequencies is outlined in section~\\ref{sc:imaging}. The effects of limited sampl\\-ing and uneven {\\it uv}--coverages are discussed in section~\\ref{sc:spsens}.  We provide analytical estimates of the sensitivity decrease, and use numerical simulations to evaluate the effect the uneven spatial sampl\\-ings have on the outcome of a comparison of VLBI images at different frequencies.  Alignment of VLBI images is reviewed in section~\\ref{sc:imalign}. A method used for spectral fitting and determining the turnover frequency is described in section~\\ref{sc:fitting}. Spectral fitting in the case of limited frequency coverage is discussed in section~\\ref{sc:frcoverage}. The first results from the turnover frequency mapping are presented in section~\\ref{sc:algorythm}. ",
        "conclusions": "}  In this paper, we have covered several methodological and scientific aspects of studying synchrotron spectrum of the parsec--scale regions in AGN. The main conclusions can be stated as follows:  1)~We have discussed a technique that can be used for mapping the turnover frequency distribution and obtaining spectral information from multi--frequency VLBA data.  A feasibility study shows that multi--frequency VLBA observations can be used for spectral imaging and continuous spectral fitting.  2)~Multi--frequency VLBA observations made with up to 10 minute separations between the scans at each frequency can provide a satisfactory spatial sampl\\-ing and image sensitivity for sufficiently bright sources with intermediate ($\\sim 10$--15\\,mas) structures. The fractional errors from comparing the data at different frequencies should not exceed 10\\% for emission with SNR$\\ge 7$, in this case.  4)~A procedure for broadband synchrotron spectrum fitting has been introduced for mapping the distribution of spectral parameters of radio emission from parsec--scale jets. Corrections based on the local curvature of the fitted spectra are introduced, in order to compensate for the incomplete frequency coverage in cases where the true turnover frequency is outside of the range of observing frequencies.  5)~From a 4--frequency VLBA observation of \\object{3C\\,345}, the first map of the turnover frequency distribution are produced. The maps indicate possible locations of the relativistic channel and strong shock fronts inside the jet. The magnetic field distribution derived from the turnover frequency and flux distributions is consistent with the plane shocks existing in the immediate vicinity of the source core. The extended emission appears to have a very low turnover frequency for which the existing data do not warrant a good estimate, limiting the conclusions to deducing certain information from the gradients of the turnover frequency which are visible in the extended jet. The observed gradients are consistent with the patterns of velocity distribution and density gradients typical for Kelvin--Helmholtz instabilities propagating in a relativistic jet. A more detailed study, with observations made at lower frequencies, is required for making conclusive statements about the nature of the observed gradients of the turnover frequency."
    },
    "9804/astro-ph9804262_arXiv.txt": {
        "abstract": " ",
        "introduction": "The EGRET experiment on board CGRO revealed the existence of a diffuse gamma-ray background (hereafter DGRB) at the level of $I_{DGRB} = 9.6 \\cdot 10^{-7} E_{GeV}^{-2.11 \\pm 0.05} ~ cm^{-2} s^{-1} sr^{-1} GeV^{-1}$ \\cite{owz94} in the energy range $0.03 \\div 10$ GeV. However, a recent reanalysis of the EGRET data \\cite{swz} found that the level of the DGRB is systematically lower by a factor $\\sim 20 \\%$ in the energy range $\\sim 0.1 \\div 4$ GeV. The DGRB is observed at high galactic latitudes $b > 10$ deg and such an evidence suggested an extragalactic origin for this diffuse background. Nonetheless, the specific origin of the DGRB is still under debate. In fact, the EGRET experiment \\cite{kan} has a poor angular resolution ($\\theta_{min} \\sim 1$ deg) so that it is hard to discriminate among different origins of this extragalactic background. Specifically, it is still difficult to discriminate between a purely diffuse nature of the DGRB (see \\eg \\cite{ca}) and the option of a DGRB  made by a superposition of unresolved, discrete sources.  The large number of identified AGNs and flat spectrum radio quasars (hereafter FSRQ) in the EGRET sky (\\cite{fi96} \\cite{has} \\cite{mat}) suggested that most of the DGRB can be produced by a non-resolved population of AGNs, the actual fraction of the DGRB produced by FSRQ and BLLacs being in the range $\\sim 40  \\div  95 \\%$ (see \\eg \\cite{com}). Separately, it has been evaluated that a fraction $\\sim 42 \\div 97 \\%$ of the DGRB could be ascribed to blazars \\cite{pad}. However, the flatness of the spectrum of the DGRB seems to favour the possibility that BLLacs could be the major contributors to the DGRB of extragalactic origin \\cite{pohl}. Erlykin \\ea \\cite{erl96} reviewed the various AGN contributions and quoted that the fraction of the DGRB produced by the observed AGNs is $\\sim 65 \\%$. The DGRB fractions previously reported may be subject to a revision ($\\sim 25 \\%$ increase) if the recent reanalysis \\cite{swz} of the EGRET data is adopted.  On account of the large theoretical uncertainties and of the present observational precision of the EGRET detectors, it is still hard to discriminate among the different proposed possibilities, even though a fluctuation analysis of the EGRET data should give more precise indications on the nature and origin of the DGRB.   Beside the discrete, unresolved source case pictured for the origin of the DGRB, there have been some pioneering works \\cite{hw} \\cite{ds} \\cite{bbp} \\cite{volk} suggesting that a relevant fraction of the DGRB could be produced by {\\it extended} sources through hadronic collisions of cosmic ray (hereafter CR) protons interacting with the protons of the Inter Galactic Medium (hereafter IGM) which is abundantly present within galaxy clusters (see \\cite{sa88} for a review). In this alternative picture, the CR's are assumed to be produced within clusters (we will discuss in Sect.4 some of the possible sources) where also a population of protons and electrons is residing in the form of a hot (with temperatures $T \\sim 10^7 \\div 10^8$ K), tenuous (with electron number densities $n_e \\sim 10^{-3}~cm^{-3}$), chemically enriched and massive (with mass fractions $M_{IGM}/M \\sim 0.05 \\div 0.3$) plasma: the IGM.  The proposed mechanism has an essential ingredient in the confinement of the CR's within clusters where they are produced; this point, already realized by some authors \\cite{bbp} \\cite{volk}, is responsible for the net increase in the probability of interaction per proton with respect to the case of a straight line propagation. The increase factor can be estimated to be $\\sim c t_{cl} /R_{cl}\\simgt 600$, where $t_{cl} \\simlt H_0^{-1}$ is the age of the cluster and $R_{cl}$ is its size. Cosmic rays produced within a cluster during all its lifetime can thus produce gamma rays through the production and the subsequent decay of neutral pions: \\begin{equation} p+p\\to \\pi^0+X~,~~~~~~~~~~~~~~\\pi^0\\to \\gamma+\\gamma. \\end{equation} Note that in the same interactions, charged pions are also produced, which determine a neutrino emission through the following channels: \\begin{equation} p+p\\to\\pi^{\\pm}+X, ~~~~~\\pi^{\\pm}\\to \\mu^{\\pm} \\nu_{\\mu}(\\bar{\\nu}_{\\mu}), ~~~~~\\mu^{\\pm}\\to e^{\\pm} + \\bar{\\nu}_{\\mu}(\\nu_{\\mu}) + \\nu_e (\\bar{\\nu}_e) ~. \\end{equation} We will also discuss the relevance of these last processes in Section 7 below.  Using the gamma ray production from clusters of galaxies according to eq. (1), Houston \\ea \\cite{hw} suggested that the total extragalactic gamma ray intensity detected above $35$ MeV \\cite{ft82}, $I_{\\gamma} \\approx 5.5 \\cdot 10^{-5}~ cm^{-2} s^{-1} sr^{-1}$, could be ascribed, for a large fraction, to galaxy clusters. They predicted a level $I_{\\gamma} \\approx 5 \\cdot 10^{-5} cm^{-2} s^{-1} sr^{-1}$ at energies above $35$ MeV, assuming an observed local cluster space density, $n_{cl} \\approx 7.3 \\cdot 10^{-5} Mpc^{-3}$, integrated out to the Hubble radius, $R_H=6 \\cdot 10^3$ Mpc, and neglecting any cosmological effect. More recently,  Dar \\& Shaviv (hereafter DS \\cite{ds}) reanalyzed the problem in the light of the EGRET data \\cite{owz94} and calculated the contribution to the DGRB from CR interactions in the intracluster gas, under the assumption that the energy density of CR's in clusters is the same as in our own galaxy (universality). With this assumption, Dar \\& Shaviv \\cite{ds} predicted a level $I_{\\gamma}(> 100~ MeV) \\approx 1.2 \\cdot 10^{-5}$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$, which could explain the whole amount of the DGRB of extragalactic origin. In a following paper, Berezinsky, Blasi \\& Ptuskin (hereafter BBP \\cite{bbp}) relaxed the {\\it ad hoc} assumption of universality, and estimated the CR energy density in clusters due to various possible sources of CR, using the condition of diffusive confinement of CRs. In their approach, BBP \\cite{bbp} showed that it is impossible to fulfill the universality condition with the usual CR sources in clusters, emphasizing that the DGRB due to the CR interactions in clusters should be a small fraction of the total diffuse flux observed by EGRET. This conclusion was reached by the previous authors under the hypothesis that a large fraction of the baryons in the universe is contained inside clusters of galaxies (BBP considered  that clusters are a fair sample of the baryons in the universe \\cite{wf91}, \\cite{wetal93}, \\cite{wf95}) assumed to have a homogeneous inner distribution of gas, $n_e =const$ (here $n_e$ is the IGM electron number density). Because of these assumptions, their results depend only on overall cosmological parameters like the baryon fraction in the universe $\\Omega_b$, and on the cluster size.  Dar \\& Shaviv \\cite{ds} also predicted the gamma ray fluxes from a few nearby  clusters (Coma, Perseus and Virgo): for these three clusters they found $\\gamma$-ray fluxes in the range $F_{\\gamma}(>100 ~MeV) \\approx 5 \\div 20 \\cdot 10^{-8} cm^{-2} s^{-1}$. In particular, the value which they predicted for A1656 (Coma), $F_{\\gamma}(>100 ~MeV) \\approx 5 \\cdot 10^{-8} cm^{-2} s^{-1}$, is close to - or slightly higher than - the upper limits given by EGRET for this source. Similar results were obtained for these clusters by Ensslin \\ea  \\cite{ensslin96} assuming a population of CRs from radio sources located within galaxy clusters in almost equipartition with the IGM thermal energy.  We stress here that in all the previous works a uniform IGM density profile was assumed. Moreover, the cluster population was not assumed to evolve with cosmic time, and the same working hypothesis of no-evolution was assumed for the IGM content of each cluster.  However, X-ray studies of galaxy clusters, have shown that these cosmic structures are indeed well structured, having a gas density profile $n(r) \\propto [1+(r/r_c)^2]^{- 3 \\beta /2}$, with core radii $r_c \\approx 0.1 \\div 0.3 \\hmpc$ and $\\beta \\approx 0.6 \\div 0.8$ (see \\eg \\cite{jf92}; see also \\cite{sa88} and references therein). Beside this, the IGM is indeed evolving as indicated by its sensitive metal enrichment, $Fe/H \\sim 0.2 \\div 0.5$ (in solar units, see \\eg \\cite{e90}, \\cite{ar}), shown even for the brighter clusters observable at redshifts $z \\sim 0.5$ \\cite{lm97}. Nonetheless, there is also an increasing debate on the possible evolution of the X-ray luminosity function observed out to $z \\simlt 0.5$ with the EINSTEIN \\cite{gio90} \\cite{h92} and ROSAT satellites \\cite{ebe97} \\cite{nich} and on the possible evolution of the cluster temperature function \\cite{cmv} \\cite{eke98} \\cite {vl98}. If an evolution is present in the cluster population this can be, in fact, understood as a result of two competing effects: \\newline {\\it i)} a luminosity evolution, where the cluster X-ray luminosity, $L \\propto n^2 T^{1/2} R^3$ (mainly due to thermal bremsstrahlung), changes with redshift due to variations in the gas mass density, $n \\propto f_g \\rho_{cl}$ (where the cluster gas mass is taken to be a fraction $f_g \\equiv M_{gas}/M$ of the total cluster mass), and/or changes in the IGM temperature $T$ at fixed mass, $M \\propto \\rho_{cl} R^3$ (here $\\rho_{cl}$ is the cluster total mass density); \\newline {\\it ii)} a change in redshift of the number density, $N(M,z)$, (usually referred to as mass function, hereafter MF) of clusters that are found to be collapsed (or virialized) in the mass range $M, M+dM$ at redshift $z$.  Detailed studies of cluster evolution in X-rays (see \\eg \\cite{cv}, \\cite{ob}, \\cite{cmv}) considered in fact that a combination of the previous mechanisms is responsible for the actual cluster evolution when they fit the available data (see \\cite{cv} for a detailed discussion).  In this paper we predict the amount of high energy, non-thermal, gamma-ray emission from galaxy clusters using detailed modelling of the realistic cluster structure, as well as viable modelling for the evolution of the IGM and of the cluster MF. Based on these phenomenological cluster models, we predict the amount of DGRB that can be produced in the viable cosmological models: here we consider flat and low-density (open or vacuum-dominated) CDM models as well as mixed Dark Matter models with a fraction $\\Omega_{\\nu} \\approx 0.3$ of the total density of the universe in form of massive neutrinos. We use $h=H_0/100$ km s$^{-1}$ Mpc$^{-1}$ throughout the paper unless otherwise specified.  The plan of the paper is the following. In Sect.2 we briefly summarize the cluster formation hystory in hierarchical scenarios for structure formation. In Sect.3 we describe a model for the production of diffuse gamma-ray emission due to the interaction of CR's with the target protons present in the extended, diffuse IGM. We consider in Sect. 4 different CR sources that can be found in connection with galaxy clusters. We discuss in Sect.5 the correlation between the extended gamma-ray emission and the much better known thermal X-ray emission coming from the IGM. Based on these properties, we construct a list of predicted $\\gamma$-ray fluxes for a compilation of X-ray clusters with detailed informations on their IGM structure, IGM temperatures  and X-ray fluxes. In Sect. 6 we present predictions for the amount of DGRB produced by galaxy clusters in different cosmological scenarios. We briefly discuss in Sect.7 the extended neutrino fluxes emerging from these objects and their contribution to a possibly detectable diffuse neutrino background (hereafter DNB). Finally, in Sect.8 we discuss our results in the light of the current limits obtained from EGRET and in the light of the future experiments for gamma-ray and neutrino astronomy. ",
        "conclusions": "In this paper we presented a detailed study of the diffuse emission of $\\gamma$-rays and neutrinos from clusters of galaxies. Using realistic modelling of the cluster structure, of their formation history and of their evolution with cosmic time, we found that galaxy clusters can provide $\\simlt 1\\%$ of the DGRB measured by EGRET (in the first release by OWZ \\cite{owz94}).  Our estimate of $I_{\\gamma}$ is quite independent on the geometry of the universe, on the assumed cosmological model and on the amount of IGM evolution, because most of their contribution to the DGRB comes from nearby, $z \\simlt 0.2$, clusters. In fact, at these redshifts the effects of curvature do not take place strongly in changing the perturbation growth factor, ${\\cal D}(z,\\Omega_0)$, (normalized at the present epoch), the difference in cluster evolution are small when the different models are normalized to the local abundance of clusters observed in X-rays and the available amount of IGM evolution - even if considered to be quite strong, $f_g \\propto (1+z)^{-1 \\div -2}$ - can provide only small variations to the cluster $\\gamma$-ray luminosities, as $L_{\\gamma} \\propto f_g$ (see eq.18). On account of all these aspects, we consider that our results for the contribution of galaxy clusters to the DGRB are quite robust.  Our approach  differs substantially from the previous ones in several (among others) aspects: \\par\\noindent {\\it i)} we considered - differently from all the previous approaches - a self-consistent approach to the formation of clusters following the spherical collapse model \\cite{peeb80} complemented with a realistic IGM density profile, consistent with the most recent determinations from X-ray observations. This fact has important effects on the CR confinement within cluster cores and hence on the relative $\\gamma$-ray and neutrino emission rates; \\par\\noindent {\\it ii)} we considered (as BBP did) here an energy dependent diffusion coefficient which results in a very general picture of the CR confinement within cluster cores; \\par\\noindent {\\it iii)} we also considered here - at variance with the previous approaches - the effects of a possible evolution in the cluster IGM content. This is consistent with the present indications of a variation in the IGM content from groups to rich clusters in the local frame and with the X-ray, shock (or entropy) induced, luminosity evolution observed from numerical simulations \\cite{tm97} and predicted in analytical models (both shock and entropy models) for the evolution of X-ray clusters \\cite{cola97} \\cite{b97}; \\par\\noindent {\\it iv)} we use the PS cluster MF that was found to be consistent with N-body simulations over a large dynamical range and up to $z \\simgt 2 $ \\cite{eke96}. We normalized it to the local abundance of clusters detected in X-rays. In the previous approaches average values for the overall cluster abundance, $n_{cl} \\approx 4\\div 7 \\cdot 10^{-5}~ Mpc^{-3}$, were used \\cite{hw} \\cite{bbp} without considering the effect of an evolving cluster mass function; \\par\\noindent {\\it v)} using a self-consistent modelling of the IGM we found an analytical correlation between $\\gamma$-ray and $X$-ray emission for clusters. The predicted ratio $F_{\\gamma} / F_X \\propto f_g^{-1} r_c^{-1} T^{-1/2}$  [see eq.(32)] provides a behaviour of the $F_{\\gamma}-F_X$ relation different from that obtained by Ensslin \\ea \\cite{ensslin96}, $F_{\\gamma} / F_X \\propto T^{1/2}$, because we did not assume any (partial) equipartition between IGM thermal energy and relativistic jet particles. Our correlation results only from the basic electromagnetic and hadronic emission mechanisms in which the IGM protons are the targets for both the X-ray thermal bremsstrahlung emission and for the $p p$ collisions responsible for $\\gamma$-rays. \\par\\noindent {\\it vi)} using such a correlation we derived a sample of nearby clusters  with predicted $\\gamma$-ray fluxes observable with the next generation $\\gamma$-ray telescopes. Incidentally, we found a $\\gamma$-ray flux for Coma, $F_{\\gamma}(>100 MeV)\\approx 8.5 \\cdot 10^{-9}~ photons$ $s^{-1} ~cm^{-2}$ which is consistent with the EGRET upper limits for this cluster (previous specific predictions \\cite{ds} \\cite{ensslin96} seem to exceed the EGRET upper limit).  Our numerical results for $I_{\\gamma,\\nu}$ are in reasonable agreement with those obtained by BBP, even though based on a quite different description of the cluster structure and evolution. This agreement is due to the fact that BBP considered a constant comoving density of clusters, $n_{cl} \\sim 5 \\cdot 10^{-5}$ Mpc$^{-3}$, assumed to be  a fair sample of the baryons in the universe, and containing a fraction $\\Omega_b\\approx 0.5 \\Omega_{BBN}$ (where $\\Omega_{BBN}$ is the value of the baryon density predicted by Big Bang Nucleosynthesis). Under these assumptions, BBP obtained a value for $I_{\\gamma}$ higher by a factor $\\sim 3\\div 4$ with respect to our result, based on values $f_g \\sim 0.1$ (see Sect.2). Our refined calculations show why their assumption of considering $n_{cl} \\sim const$ was reasonable:  the clusters effectively contributing to the DGRB are located at $z \\simlt 0.2$ (see Fig.5), where the effects of evolution do not have room to take definitely place (see Fig.7).   Because of the inherent uncertainties in the predictions of quantities whose calculation involve to set the values of parameters which are not known precisely, we also estimated the range spanned by $I_{\\gamma, \\nu}$ for the combination of parameters allowed by the present observational ranges. In fact, the description of the cluster structure and evolution that we used in our analytic approach consider only ensemble averaged quantities. But we observe a whole distribution of the real cluster properties with respect to the average cluster moulding. Some amount of variance is needed to be considered in cluster modelling in order to ensure the predictive power of the viable models for structure formation. To explore the role of the uncertainties in the relevant quantities  we considered several sources of uncertainties (see Sect.7).   From an inspection of Figs. 9 and 10 we note that the effects of  the possible theoretical uncertainties in the description of the cluster and IGM evolution could change the predicted contributions for the DGRB and for the DNB by a factor $\\simlt 3$, setting the maximal level of $I_{\\gamma}$ to a few $\\%$ of the EGRET value.   The DGRB seems to be mostly produced by AGNs (FSRQ  and/or BLLacs) and/or blazars (we take here an  estimate of $\\sim 60 \\div 65 \\%$ \\cite{erl96} of the EGRET diffuse flux \\cite{owz94}). Diffuse $\\gamma$-ray emission could be observed also from normal galaxies yielding a contribution $\\sim 5 \\%$ \\cite{erl96}. When added to the $\\sim 10 \\div 15 \\%$ of the DGRB contributed by their high-$E$ photons interacting with other existing backgrounds (\\eg IR, CMB \\cite{erl96}) one gets  only $\\sim 15 \\div 20 \\%$ of the DGRB left for truly diffuse or extended sources. Of this amount, a fraction of the diffuse $\\gamma$-ray flux $\\sim 3 \\div 5 \\%$ is predicted \\cite{ww} to originate from decaying topological defects (see \\cite{sigl}) and interactions of UHE particles with the CMB. Note, however, that the amount and the spectral distribution of this possible diffuse background depend sensitively on the amplitude of the primordial magnetic field on scales larger than supercluster sizes. So, according to the previous estimates, the presence of all these sources of diffuse $\\gamma$-ray emission (even though partially model dependent) determines an upper limit to the contribution of extended extragalactic sources to the DGRB, that is $\\sim 10 \\div 22 \\%$ of the OWZ EGRET level \\cite{owz94}. This sets  rather weak constraints on the level of CR production in clusters and hence on the presence and activity of AGNs in clusters or on the formation and efficiency of accretion shocks around clusters. However, if we consider the revised level of the DGRB as derived by SWZ \\cite{swz}, then the previous upper limit reduces to $\\simlt 2 \\%$. Our predictions of the DGRB contributed by galaxy clusters $I_{\\gamma,cl} / I_{EGRET} \\sim 0.005 \\div 0.02$ is perfectly compatible with the presence of both a population of evolving FSRQ and AGNs dominating the $\\gamma$-ray sky and with the presence of truly diffuse backgrounds like those previously discussed. Note, however, that the major source of uncertainty in the level of the extragalactic DGRB comes from the contribution of the AGNs. A fluctuation analysis of the EGRET data is needed to have more definite indications on the level of the DGRB contributed by discrete sources. If, on the other hand, CR acceleration will be found to be relevant in clusters (yielding $L_{CR}$ substantially  larger than $ 10^{44}$ erg/s), then the predicted level of DGRB produced by galaxy clusters can set interesting limits to the space density and evolution of $\\gamma$-ray AGNs.   The sensitivities and angular resolutions achievable by the next generation gamma-ray (INTEGRAL, GLAST, AMS) and neutrino (see \\cite{halz} for a review) detectors will be able to shed a new light on the high energy phenomena occurring in large scale structures.   \\vskip 1.5truecm \\newline {\\bf Acknowledgements} We thank the Referee for useful comments and suggestions which improved substantially the presentation of the paper. We also aknowledge interesting and stimulating discussions with V.S. Berezinsky during a recent visit of S.C. at the LNGS. S.C. aknowledges also interesting discussions with G. Kanbach, A. Dar and F. Halzen, among others, at the 1997 Moriond Meeting {\\it High Energy Phenomena in the Universe}. The research of P.B. is funded by a INFN PostDoctoral Fellowship at the University of Chicago.  \\newpage"
    },
    "9804/astro-ph9804054_arXiv.txt": {
        "abstract": "We present the results of an {\\asca} observation of the LINER NGC~4579. A point-like X-ray source is detected at the nucleus with a 2--10~keV luminosity of $1.5\\times10^{41}$~{\\eps}assuming a distance of 16.8~Mpc. The X-ray spectrum is represented by a combination of a power-law with a photon index of $\\sim$1.7 and soft thermal component with $kT\\sim$0.9~keV. An iron K emission line is detected at $6.73\\pm0.13$~keV (rest frame) with an equivalent width of $490^{+180}_{-190}$~eV and is statistically significant at more than 99.9\\% confidence. The line center energy is consistent with Helium-like iron and is significantly higher than 6.4~keV which is expected from fluorescence by \"cold\" (or a lower ionization state of) iron. The iron line profile shows no significant red tail in contrast to Seyfert 1 galaxies although the statistics are limited. The line center energy, equivalent width, and profile are consistent with an origin in an ionized accretion disk. However the large mass accretion rate necessary to ionize the accretion disk is not consistent with the observed luminosity and normal accretion models. ",
        "introduction": "Recent optical spectroscopic surveys have shown that there are many active galactic nuclei (AGNs) in nearby galaxies and about 40\\% of bright galaxies are classified as Seyfert galaxies or LINERs (Low Ionization Nuclear Emission-line Regions; \\cite{hec80}) (Ho, Filippenko, \\& Sargent 1997a). The luminosity of these objects are rather low compared to previously known AGN with a median value of the H$\\alpha$ luminosity being only $2\\times10^{39}$ {\\eps} in the sample of Ho et al (1997a). Such objects (low luminosity AGNs; hereafter LLAGNs) are important for investigating the physics of AGN under an extreme condition, i.e. very low luminosity. X-ray observations probe the innermost regions of AGNs and specifically the iron K line provides information on the ionization state, density, and motion of matter very close to the central energy source.  {\\asca} observations of Seyfert 1 galaxies revealed that as a class these objects have a broad iron K line with a profile skewed to lower energies, thought to be caused by the reprocessing of the continuum by a relativistic accretion disk (e.g., \\cite{tan95}, Nandra et al. 1997a). The center energy of the iron line from Seyfert 1 galaxies is consistent with 6.4~keV, which is expected from fluorescence by neutral or lower ionization states of ($<$\\ion{Fe}{16}) iron in a disk with an inclination of $<$ 30 degrees. In the Seyfert 1.9 galaxy IRAS ~18325--5926, a higher peak energy of iron emission is seen, which is compatible with a highly-inclined disk ($i=40-50^{\\circ}$) origin (\\cite{iwa96}). Highly-ionized iron emission lines are detected from several radio-quiet quasars e.g. E1821+643 (\\cite{kii91}, \\cite{yam97}) and PG~1116+215 (\\cite{nan96}). Nandra et al. (1997b) studied the luminosity dependence of the iron line profile in a large sample of AGN and found that the center energy increases and the red-tail becomes weaker with increasing luminosity. They attributed such behavior to an increasing ionization of the accretion disk with increasing luminosity. Thus X-ray measurements of iron emission lines are powerful diagnostic tools of matter in the vicinity of the nucleus.  There are only a few observations of iron emission lines from low luminosity AGNs ( {\\LX} (2--10~keV) $\\sim 10^{40}-10^{41}$~{\\eps}). M81 (NGC~3031) with an X-ray luminosity of {\\LX} (2--10~keV) $\\sim 2\\times10^{40}$~{\\eps} shows a broad iron line centered at $\\sim6.7$~keV with an equivalent width of $\\sim 200$~eV. This line center energy is significantly higher than Seyfert 1 galaxies and similar to luminous quasars. An iron line at 6.4~keV with an equivalent width of $\\sim300$~eV is detected from the low luminosity Seyfert 1 galaxy NGC~5033 ({\\LX} (2--10~keV) = $2\\times10^{41}$~{\\eps}, \\cite{te98b}) but only an upper limit on the equivalent width of {\\simlt} 300~eV is obtained for NGC~1097 ({\\LX} (2--10~keV) = $1\\times10^{41}$~{\\eps}, Iyomoto et al. 1996). Although strong iron emission lines are also detected from M51 (= NGC~5194, Terashima et al. 1998a), NGC~1365 and NGC~1386 (Iyomoto et al. 1997), the iron lines in these objects are interpreted as being caused by reprocessed emission from an obscuring tori and/or extended ionized scatterer outside of our line of sight, i.e., these nuclei are heavily obscured. Thus, at present, the number of LLAGNs with small intrinsic absorption from which iron lines are detected is rather limited.  NGC~4579 (M58) is a Sab galaxy in the Virgo cluster of galaxies and classified as a LINER or Seyfert 1.9 galaxy based on the optical emission lines (\\cite{ho97a}, \\cite{kee83}, \\cite{sta82}) and the broad H$\\alpha$ component, detected with a FWHM $\\sim 2300$~km s$^{-1}$ (\\cite{ho97b}). There exists a flat-spectrum radio core (\\cite{hum87}). An {\\Einstein} HRI observation showed the presence of an unresolved X-ray nucleus and the X-ray flux was measured to be {\\FX} = $7.9\\times10^{-12}$~{\\eps} cm$^{-1}$ in the 0.2--4.0~keV band with the {\\Einstein} IPC (\\cite{fab92}, Halpern \\& Steiner 1983) which corresponds to the X-ray luminosity of $2.7\\times10^{41}$~{\\eps} (we assume a distance of 16.8~Mpc (\\cite{tul88}) throughout this paper). These facts indicate the presence of a LLAGN in this galaxy. A recent ultraviolet imaging observation by {\\it Hubble Space Telescope} ({\\HST}) Faint Object Camera (FOC) detected a point source at the nucleus (\\cite{mao95}). Its UV spectra were taken by {\\HST} Faint Object Spectrograph (FOS) and a featureless UV continuum is detected as well as various emission lines. Comparison of the FOC and FOS data also indicate a factor of 3.3 decrease of UV flux in 19 months. The narrow UV emission lines are incompatible with shock excitation model and a photoionization model is preferred (Barth et al. 1996). Several broad UV emission lines are also detected. These UV results provide further support for the presence of a LLAGN in NGC~4579. On the other hand, \\cite{mao98} estimated the ionizing photon number by extrapolating the UV luminosity at 1300 A towards higher energies and argued that the observed UV continuum is not sufficient to explain the H$\\alpha$ luminosity. They also suggest that emission from AGNs is most prominent at energies higher than the UV. Measurements of an X-ray flux and continuum slope provide information on the ionization source in this LINER.  In this paper we report the detection of an Iron K emission line centered at 6.7~keV and discuss X-ray properties of the LLAGN in NGC~4579 and origin of the iron emission line. ",
        "conclusions": "\\subsection{X-ray emission from a low luminosity AGN}  We obtained X-ray images and spectra in the 0.5--10~keV band and a point-like X-ray source with a photon index of $\\Gamma = 1.72\\pm0.05$ is detected. An iron line is also detected at 6.7~keV. In the soft energy band, a broad line like feature identified with iron-L line complex indicates the presence of thin-thermal plasmas of temperature $kT\\sim 0.9$~keV.  The X-ray luminosity ($1.5\\times 10^{41}~${\\eps} in 2--10~keV ) is 1--3 orders of magnitude smaller than typical Seyfert galaxies and falls in the classes of LINERs and \"low luminosity\" Seyfert galaxies (Serlemitsos et al. 1996, Iyomoto et al. 1996, Ishisaki et al. 1996, Terashima et al. 1998b). In normal spiral galaxies, the X-ray emission is dominated by discrete sources, specifically low mass X-ray binaries (LMXBs) (e.g. \\cite{fa89}, Makishima et al. 1989). The X-ray luminosity from LMXBs are roughly proportional to B-band luminosity {\\LB} and their X-ray spectra can be approximated by a thermal bremsstrahlung of a temperature of several keV. The {\\asca} X-ray spectrum of NGC~4579 is also fitted by $kT\\sim8$~keV thermal plasma model. However the strong iron line at 6.7~keV is not compatible with the X-ray spectra of LMXBs, since the equivalent width of iron emission lines from LMXBs are small (several tens of eV, \\cite{hir87}). Additionally, the {\\LX}/{\\LB} value $1.3\\times10^{-3}$ is more than an order of magnitude higher than normal spiral galaxies (e.g. \\LX /\\LB =$3.5\\times10^{-5}$ for M31; \\cite{mak89}). Additionally an upper limit on the size of an archival {\\rosat} PSPC image is 14\" (Gaussian $\\sigma$), which corresponds to 1.1 kpc at 16.8 Mpc. This upper limit is significantly lower than the size of the galaxy.  Therefore we conclude that contribution from LMXBs to X-ray emission of NGC~4579 is negligible.   Hot plasmas with temperatures on the order of $\\sim10$ keV are present in the Galactic center region and their X-ray spectra show prominent, ionized iron K emission (e.g. Koyama et al. 1996).  The X-ray spectral shape of NGC 4579 in the hard X-ray band is similar to such hot gas. However, the X-ray luminosity of NGC 4579 is three orders of magnitude higher than the Galactic ridge emission ({\\LX}$\\sim2\\times10^{38}$ {\\eps}; Kaneda et al. 1997, Yamasaki 1996, Warwick et al. 1985).  Starburst galaxies also show a hard spectral component with a temperature of $\\sim$10 keV and their X-ray luminosities are around $10^{40}$ {\\eps} (e.g. $3.4\\times10^{40}$ \\eps in 2--10 keV for M82; Ptak et al. 1997). However the starburst activity in NGC 4579 is weaker than M82, since the far-infrared luminosity of NGC 4579 is about an order of magnitude lower than that of M82. Furthermore, starburst galaxies show weak or no iron emission contrary to NGC 4579. Therefore hot plasma is unlikely as the origin of the hard component and iron emission line in NGC 4579 and we conclude that the AGN emission dominates the {\\it ASCA} spectra and that other components such as a hot gas contribution is small, if any.  We note that errors in background subtraction of the Virgo cluster hot gas do not affect the detection of the iron emission line at 6.7 keV, since the cluster gas is very dim in this region and temperature is low ($kT\\sim2$ keV; Matsumoto 1998, \\cite{boh94}). Actually no significant iron emission is detected from the GIS field around NGC 4579.   If the primary ionizing mechanism of LINER optical emission lines in this galaxy is photoionization by a LLAGN, {\\LX}/{\\LHa} might be expected to be similar to Seyfert 1 galaxies, for which there is a good positive correlation between {\\LX} and {\\LHa} (e.g. Ward et al. 1988, \\cite{kor95}, Serlemitsos et al. 1996). Using the H$\\alpha$ luminosity of broad plus narrow component {\\LHa} = $5.9\\times10^{39}$~{\\eps} (\\cite{ho97b}) and the observed X-ray luminosity in the 2--10~keV band, we obtain {\\LX}/{\\LHa} $\\approx$ 26 for NGC~4579. This value is in excellent agreement with those of Seyfert 1 galaxies (\\cite{war88}) and strongly supports a low luminosity AGN as the ionizing source of the LINER in NGC~4579.  Less luminous Seyfert 1 galaxies tend to show rapid and large amplitude variability (Mushotzky, Done, \\& Pounds 1993 and references therein). However NGC 4579 shows no significant short term variability. Lack of variability on short time scales seems to be a common property of LLAGNs (\\cite{mus92}, \\cite{pet93}), for example the LLAGN in NGC 1097 (Iyomoto et al. 1996) and NGC 3998 (\\cite{awa92}) also show no significant variability on timescales less than a day. Direct comparison of {\\rosat} PSPC and {\\asca} flux in the 0.5--2 keV band show a factor of two increase in $\\sim$3.5 years.  The X-ray spectral slope $\\Gamma = 1.72\\pm0.05$ is identical to the average value found for hard X-ray selected Seyfert 1 galaxies (Mushotzky et al. 1993) but the luminosity is lower than that of any Seyfert 1 galaxy but NGC 4051. Based on the FW0I (full width at 0 intensity) of a broad emission line and an estimate of the size of the broad line region, mass of the central black hole is roughly estimated to be $M_{\\bullet} \\sim 4\\times10^6M_{\\odot}$. Then the Eddington ratio $L/L_{\\rm Edd}$ is $\\sim 10^{-3}$ for the observed luminosity of $\\sim5\\times10^{41}$~{\\eps} (Barth et al.  1996), although their blackhole mass estimation is crude. Therefore the X-ray spectral slope does not seem to be drastically changed even at a very low Eddington ratio. This is also true for M81, for which $L/L_{\\rm Edd}$ is estimated to be $\\sim (2-10)\\times 10^{-4}$ (\\cite{ho96}) and the photon index is $1.85\\pm0.04$ (\\cite{ish96}).  Soft thermal emission of $kT\\sim0.5-1$~keV is often observed from low luminosity AGNs (Terashima 1997, \\cite{pta97}, Serlemitsos et al. 1996). In some cases, such emission is associated with starburst activity (e.g. \\cite{iyo96}, \\cite{te98a}).  Since the far-infrared luminosity of NGC~4579 is $1.5\\times10^{43}$~{\\eps} some star formation activity may be present which may explain the thermal emission. The soft thermal X-ray to far infrared luminosity ratio {\\LX}/{\\LFIR} = $6\\times10^{-4} - 1.1\\times10^{-3}$ is consistent with starburst galaxies (e.g. \\cite{dav92}) within the scatter.   \\subsection{Iron-K line}  A marginally broad ($\\sigma \\approx 0.17$~keV) iron emission line is clearly detected at $6.73^{+0.13}_{-0.12}$~keV and the equivalent width is $490^{+180}_{-190}$~eV for the broad Gaussian model fit. The line center energy is significantly higher than 6.4~keV, which is typically observed from Seyfert 1 galaxies, and consistent with He-like iron.  A similar broad iron line centered at $\\sim 6.7$~keV is detected from the low luminosity Seyfert galaxy M81 (Ishisaki et al. 1996, Serlemitsos et al. 1996). The line can also be represented by line blending of neutral, He-like, and H-like iron and dominated by He-like iron.  The disk-line profile (Fabian et al. 1989) is probably inconsistent with the data for 6.4~keV or 6.7~keV intrinsic line energy because of following reasons. The $\\chi^2$ value is worse than a single broad Gaussian fit and systematic residuals remain in the disk-line fit, since a significant red tail is not clearly seen in the data.  Furthermore, the disk-line model provides the very large equivalent width $\\sim900$~eV, which is about 4 times larger than the results of the disk-line fit to Seyfert 1 galaxies ($<$EW$>=(230\\pm60)$ ~eV, \\cite{na97a}). Therefore our data prefer a symmetric Gaussian-shape profile with intrinsic line center energy of 6.7~keV (He-like) rather than 6.4~keV ($<$\\ion{Fe}{16}). Thus the ionization state of the iron line emitter may be different from that of higher luminosity Seyfert 1 galaxies in at least some LLAGNs (NGC 4579 and M81)   Strong ionized iron emission lines are observed in heavily obscured Seyfert 2 galaxies (NGC 1068, \\cite{uen94}, \\cite{iwa97}; NGC 1365, \\cite{iyo97}; see also \\cite{tu97a}, \\cite{tu97b}). In these objects continuum emission from the nucleus is completely blocked and only scattered radiation is observed.  Ionized iron lines are interpreted as originating from a photoionized scattering medium. If the continuum of NGC~4579 is scattered radiation, then the observed X-ray luminosity is only a fraction of its intrinsic luminosity. Since the scattering fraction is typically less than 10 \\% for Seyfert 2 galaxies (\\cite{uen95}), the {\\LX}/{\\LOIII} should be less than 10 \\% of those of Seyfert 1 galaxies as is the case for NGC 1068 (\\cite{mul94}).  However the observed X-ray to [OIII]$\\lambda$5007 luminosity ratio {\\LX}/{\\LOIII} is very similar to Seyfert 1 galaxies. Therefore the observed X-ray continuum is not likely to be due to a scattered component. Then the observed iron line should be emitted from the matter close to the nucleus in order to be ionized and/or broadened due to the Doppler effect.   If the iron line is emitted by an accretion disk, a line profile with significant red tail is expected (\\cite{fab89}). On the other hand, the observed profile seems to be symmetric in shape although the statistics are limited. Broad lines with weaker red tails than Seyfert 1s are observed in AGNs with much higher luminosity; {\\LX}$>10^{44}$~{\\eps} (\\cite{na97b}). If the inner-most part of the accretion disk is almost fully ionized, the red component is expected to be very weak or absent. Thus the observed profile is consistent with the interpretation that the observed iron K emission is from an ionized disk.   The obtained equivalent width ($\\sim 500$~eV for the Gaussian model) is rather large compared to that seen in most Seyfert 1 galaxies.  If the disk is highly ionized, the fluorescence yield of iron increases and absorption by lighter elements decreases as light elements are almost completely ionized.  In such a situation the equivalent width of an iron line can increase by a factor of two (Matt et al. 1993, \\cite{zyc94}). Therefore the large equivalent width is also naturally explained by an ionization effect.  The ionization state of photoionized matter is determined by an ionization parameter $\\xi = L/nR^2$ (\\cite{kal82}), where $L$, $n$, and $R$ is the luminosity of ionizing photons, the number density of photoionized matter, and the distance from light source to photoionized matter, respectively. The X-ray luminosity of NGC 4579 is only $1.5\\times10^{41}$ {\\eps}, which is 1--3 orders of magnitude smaller than for Seyfert 1 galaxies, and the X-ray luminosity of M81, from which an iron line centered at $\\sim6.7$~keV is detected, is even lower ($\\sim 2\\times10^{40}$ {\\eps}). In order to photoionize iron atoms to be He-like, $\\xi$ should be at least $\\sim$ 500, while $\\xi<100$ is required for less ionized species ($<$ \\ion{Fe}{16}) which is probably appropriate for usual Seyfert 1 galaxies. Therefore $nR^2$ in the iron line emitting region should be more than two orders of magnitude smaller than that of luminous Seyfert 1 galaxies. An expected ionization parameter under an assumption of standard $\\alpha$ disk is calculated by \\cite{mat93}. According to their results, the ionization parameter has a strong dependence on the mass accretion rate $\\xi \\propto \\dot{m}^3$ (equations (5) and (6) in Matt et al. 1993), where $\\dot{m}$ is denoted in units of the critical accretion rate $\\dot{m} = L/L_{\\rm Edd}$. In order to ionize iron to He-like, $\\dot{m}$ should be at least 0.2 (Figs. 2 and 5 in Matt et al. 1993). However the order-of-magnitude estimate of the central black hole mass by \\cite{bar96} combined with the observed luminosity gives a significantly smaller value of $\\dot{m}$ $\\sim1\\times10^{-3}$. Then we cannot explain the very low luminosity and the ionized iron line at the same time in the standard disk model. This may suggest that the accretion processes in AGN is different in a very low luminosity situations with very small $\\dot{m}$.  An advection dominated accretion flow (ADAF) model is proposed for AGNs specifically for objects radiating at very low Eddington ratio (e.g. $\\dot{m} \\sim 10^{-4}$ for NGC 4258, Lasota et al. 1996). In the model by Lasota et al. (1996), a standard disk is assumed outside of $r_{\\rm in}$ and an ADAF inside of $r_{\\rm in}$. In an ADAF, accreting matter is heated up to very high temperatures ($T_i\\sim10^{12}$K, $T_e\\sim10^{9}$K).  However our detection of an iron line indicates the presence of highly ionized (but not fully ionized)  matter surrounding a large solid angle viewed from the light source.  This means that $r_{\\rm in}$ should be small and a geometrically thin disk is appropriate. Therefore the iron line in NGC~4579 cannot be explained solely by an ADAF model and the real situation in NGC~4579 may correspond to a condition near the transition from the $\\alpha$ disk to an ADAF.  Future sophisticated modeling of accretion in LLAGNs and calculation of expected iron emission as well as precise measurements of an iron K line and mass determination by {\\HST} Space Telescope Imaging Spectrograph will be important to understand physical processes in extremely low luminosity AGNs.\\\\"
    },
    "9804/gr-qc9804048_arXiv.txt": {
        "abstract": "We conduct a direct comparison of three different representative numerical codes for constructing models of rapidly rotating neutron stars in general relativity. Our aim is to evaluate the accuracy of the codes and to investigate how the accuracy is affected by the choice of interpolation, domain of integration and equation of state. In all three codes, the same physical parameters, equations of state and interpolation method are used. We construct 25 selected models for polytropic equations of state and 22 models with realistic neutron star matter equations of state. The three codes agree  well with each other (typical agreement is better than 0.1 \\% to 0.01 \\%) for most models, except for the extreme assumption of uniform density stars.  We conclude that the codes can be used for the construction of highly accurate initial data configurations for polytropes of index $N>0.5$ (which typically correspond to realistic neutron stars), when the domain of integration includes all space and for realistic equations with no phase transitions. With the exception of the uniform density case, the obtained values of physical parameters for the models considered in this paper can be regarded as ``standard'' and we display them in detail for all models. ",
        "introduction": "The physical state of the neutron star matter has not been fully understood yet because it is very difficult to investigate particle interactions beyond nuclear matter density ($\\varepsilon_{\\rm N}/c^2 \\sim 2 \\times 10^{14}$ g cm$^{-3}$) either from nuclear experiments or from nuclear theories, (here $\\varepsilon_{\\rm N}$ is the energy density of the nuclear matter and $c$ is the velocity of light). Given this situation, one promising approach to explore the behavior of very high density matter is to make use of the macroscopic quantities of neutron stars. In particular, the mass and the rotational period of neutron stars depend crucially on the softness of the equation of state (EOS) at very high densities (see e.g. Friedman et al. 1984, 1986, 1989), thus, observational constrains, matched with theoretical models, may help in reconstructing the equation of state of very high density matter.  Given a particular equation of state, the mass of neutron stars varies with central energy density and always reaches a maximum. This implies that if the maximum mass of neutron star models constructed with a certain equation of state is smaller than the mass of observed neutron stars, that equation of state must be discarded. Currently, the largest accurately measured mass of slowly rotating neutron stars is $M_{\\rm BP} = 1.44 M_{\\odot}$, where $M_{\\rm BP}$ is the mass of one of the components of the binary pulsar PSR1913+16 (Taylor \\& Weisberg \\cite{tayl89}) and $M_{\\odot}$ is one solar mass.  Individual masses of neutron stars have also been estimated in six other binary pulsars (Thorsett et al. \\cite{thor93}, Wolszczan \\cite{wol97}), as well as in six X-ray binaries (van Kerkwijk et al. \\cite{vankerk95}) but the accuracy is not as good as in PSR1913+16. Thus, equations of state which give larger masses than $M_{\\rm BP}$ for slowly rotating stars, can be valid as candidates for the real equation of state at very high densities. Since the maximum mass of neutron stars is smaller for more compressible (soft) equations of state than for less compressible (stiff) equations of state, the true equation of state at high densities cannot be extremely soft.  On the other hand, stiff equations of state can be limited by considering the neutron star with the shortest rotational period, i.e. the most rapidly rotating pulsar.  There exists a lower limit on the rotational period for each equation of state, because if the centrifugal force exceeds the self-gravity at the equatorial surface, no equilibrium states are allowed. The lower limit of the rotational period depends on the softness of the equation of state - the radius of neutron stars with softer equations of state is smaller, which allows for higher rotation rates. Thus, if very rapidly rotating neutron stars should be found, we could exclude most stiff equations of state. At the moment, the shortest period of observed pulsars is 1.56 ms, of PSR1937+21. Consequently, equations of state for which the shortest rotational period is larger than this value, must be excluded as candidates for the real equation of state for neutron star matter.  The discussions above require us to make use of highly accurate schemes for constructing rotating neutron star models, in order to compute precise theoretical values of masses and rotational periods. Highly accurate relativistic equilibrium models are also needed as initial data for relativistic time-evolution codes (modeling of nonlinear pulsations, collapse and generation of gravitational waves). Recently, a number of groups have succeeded in constructing models of rapidly rotating neutron stars (Friedman et al. 1984, 1986, 1988, 1989, Eriguchi et al. \\cite{erig94}, Salgado et al. \\cite{salg94}, \\cite{salg94b}, Cook et al. \\cite{cook94b}, Stergioulas \\& Friedman \\cite{stag95} -- for a recent review see Stergioulas \\cite{S98}).  However, the obtained models by those authors do not always agree with each other (see e.g. Friedman et al. \\cite{frie89}, Eriguchi et al. \\cite{erig94}, Salgado et al. \\cite{salg94}, Cook et al. \\cite{cook94b}, Stergioulas \\& Friedman \\cite{stag95}). Although Stergioulas \\& Friedman~(1995) have determined the cause of the discrepancy between models in Friedman et al.~(1989) and Eriguchi et al.~(1994), (which was due to the use of a slightly different equation of state table), the reasons of smaller differences which remain, even after using exactly the same equation of state, have not been clarified yet. This is because numerical techniques used in the different codes, such as the choice of parameters defining the model, the interpolation method, the method of integrating the field equations, a.s.o. are not the same.  In this paper, three groups using their own codes (Komatsu et al. \\cite{koma89a}, Eriguchi et al. \\cite{erig94}, Salgado et al. \\cite{salg94}, Stergioulas \\& Friedman \\cite{stag95}) will decrease the differences between their results to a minimum possible, by tuning each code and using the same parameters, the same schemes of interpolation, the same equations of state, and so on.  Since the basic schemes used by the three groups are different, it will be impossible to have exactly the same results and the relative differences between results are a measure of the accuracy of the codes.  Models obtained with small relative differences between the three codes can be considered as ``standard\" models for each equation of state.  Furthermore, this direct comparison allows us to investigate the effect that the choice of interpolation method, equation of state and domain of integration has on the accuracy of the codes. ",
        "conclusions": "\\subsection{Discussion}  \\subsubsection{Metric Potentials}  As redshift factors differ by about 10 \\%  for the constant density, relativistic model N00$rr$ between the three codes (while the agreement global quantities is within a few  \\%) we compare directly the local values of metric potentials for several models.  Figures 1 to 16 show the four metric potentials (upper panel) and the relative differences in them between BGSM and KEH(OR) (middle panel) and KEH(SF) and BGSM (lower panel) for the  models N05$mr$, N15$mr$, L(L)$mr$ and WFF(FPS)$br$. The metric potentials are graphed against the coordinate $r$ in the equatorial plane ($\\theta= \\pi/2$, solid line) and along the axis of rotation ($\\theta=0$, dashed line). The range of the coordinate $r$ is the twice the equatorial radius of the star.  In general, the agreement in the local values of the metric potentials reflects the agreement in the computed physical parameters of models. In these graphs, several significant behaviors can be pointed out: First, there are high frequency and small amplitude oscillations at the inner part of the stars for all models. Second, the differences between the results of KEH(OR) and those of the other two codes are growing outside the stars as $r$ increases. Third, although the differences between the KEH(SF) and BGSM codes are very small for models N05$mr$, N15$mr$ and WFF(FPS)$br$, there appear larger differences for the stiff  model L(L)$mr$. Fourth, there appears a larger amplitude oscillation in the metric potential  $\\omega$ on the axis of rotation, close to the surface.  The first behavior is due to the integration scheme of the KEH code, i.e. the Simpson scheme. In general, the Simpson scheme gives results with higher precision, compared with those obtained by the trapezoidal scheme.  However, in the KEH scheme, the integrands contain nonsmooth functions with respect to the radial coordinate, because of the nature of the Green's functions. The non-uniform distribution of the weight factor in Simpson's scheme for nonsmooth functions results in oscillating behaviors with very small amplitudes, which cannot be noticed in the behavior the integrated quantities.  The second behavior in the original KEH code is caused by the \"truncation\" of the domain of integration at a finite distance from the star, instead of integrating over the whole space.  The large differences in the metric potentials between KEH(SF) and BGSM for EOS L, could be accounted to the stiffness of the equation of state, but the differences between KEH(OR) and BGSM for the same model are not as large, and we have not an explanation for that.  The oscillations in $\\omega$ on the axis of rotation near the surface are present also for the soft $N=1.5$ polytropes, while for $N=0.5$ they are larger. This indicates that terms in the field equations for $\\omega$ are very sensitive to the presence of the surface and the accompanying Gibbs phenomenon. Even for $N=1.5$ polytropes, where the density goes to zero smoothly at the surface, there is a small scale Gibbs phenomenon, due to the finite number of grid points used to represent the region of integration.    \\subsubsection{Method of Interpolation}  An important factor for the local accuracy of models constructed with realistic equations of state is the method of interpolation of the energy vs. pressure data given in an EOS  table. While global quantities are not affected significantly, the virial identities for realistic EOSs, are sensitive to the interpolation scheme This can be considered to reflect the nature of the interpolation scheme as mentioned before. If we define the enthalpy ($H$) by \\begin{equation} H \\equiv \\ln \\left( { \\varepsilon + p \\over \\rho c^2} \\right), \\end{equation} the Gibbs-Duhem relation, which follows directly from the first law of thermodynamics, implies \\begin{equation} {dp\\over dH} = \\varepsilon + p  \\ . \\end{equation} In the cubic Hermite interpolation, the Gibbs-Duhem relation is used to replace by $\\nabla H$ the term $\\nabla p /(\\varepsilon+p)$ which appears in the hydrostationary equilibrium equation. If the tabulated function $p(H)$ fails to satisfy the above relation, then the hydrostationary equilibrium equation, which is derived from the Bianchi identity, is only approximately verified by the numerical solution, which  results in increased error in the GRV2 and GRV3 virial identities.  The four point Lagrange interpolation does not satisfy the Gibbs-Duhem relation because it only reproduces the values of the discrete points, but there is no guarantee for the reproduction of the derivatives. This explains why the GRV2 and GRV3 errors are bad, even in the nonrotating case (GRV2 = 3E-03, GRV3 = 1E-02 for model L$sr$) as compared to  ${\\rm GRV2}\\sim 10^{-14}$ for polytropic models (see e.g. Bonazzola et al. \\cite{bona93}). The GRV2 and GRV3 error indicators thus do not reflect the precision of the code but the bad thermodynamical behavior of the tabulated EOS.  The advantage of the cubic Hermite interpolation is that the Gibbs-Duhem relation is automatically satisfied because this interpolation reproduces not only the values themselves but also the derivatives (Swesty 1996). Moreover, in our case, the energy density and the baryon number density are obtained by \\begin{eqnarray} \\varepsilon & = & {p \\over H} {d\\log p\\over d\\log H} - p,  \\\\ n & = & {\\varepsilon + p \\over m_{\\rm B} c^2} \\exp(-H) \\ . \\end{eqnarray} Because of these equations, the Gibbs-Duhem relation is satisfied in every point. Note also that the value of $\\varepsilon$ obtained in this way coincides exactly with $\\varepsilon_i$ at the points in the tabulated equation of state.    \\subsection{Conclusion}  The comparison of three different codes for constructing rapidly rotating relativistic neutron star models demonstrates that the BGSM and KEH schemes used are highly accurate for typical polytropic models - when the field equations are solved to infinity - and for models constructed with realistic equations of state, that do not have phase transitions.  If one approximates neutron stars as constant density stars, then Gibbs phenomena at the discontinuous surface reduce the accuracy of the computed models.  If high accuracy in such models and in models with phase transitions is desired, then modified numerical schemes - free of Gibbs phenomena - need to be used.  Such numerical schemes could employ, for example, surface fitted coordinates. Such a scheme has been presented recently by Bonazzola et al. (1998a) in the framework of spectral methods and looks promising for rotating stellar models.  Further, we demonstrated that the metric potentials are subject to various local oscillatory behaviors, even if integrated quantities have very good accuracy.  This observation is important for the effort of constructing initial data for the numerical evolution of rotating relativistic neutron star models.    \\begin{figure} \\resizebox{\\hsize}{10.5cm}{\\includegraphics{zeta_19.eps}} \\caption{Same as Figure 1 but for the metric potential $\\zeta$ of model N15$mr$.} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{10.5cm}{\\includegraphics{zeta_27.eps}} \\caption{Same as Figure 1 but for the metric potential $\\zeta$ of model WFF(FPS)$br$.} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{10.5cm}{\\includegraphics{zeta_4.eps}} \\caption{Same as Figure 1 but for the metric potential $\\zeta$ of model N05$mr$.} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{10.5cm}{\\includegraphics{zeta_26.eps}} \\caption{Same as Figure 1 but for the metric potential $\\zeta$ of model L(L)$mr$.} \\end{figure}"
    },
    "9804/gr-qc9804081_arXiv.txt": {
        "abstract": "In this paper we investigate extended inflation with an exponential potential $V(\\sigma)= V_0\\,e^{-\\kappa\\sigma}$, which provides a simple cosmological scenario where the distribution of the constants of Nature is mostly determined by $\\kappa$. In particular, we show that this theory predicts a uniform distribution for the Planck mass at the end of inflation, for the entire ensemble of universes that undergo stochastic inflation. Eternal inflation takes place in this scenario for a broad family of initial conditions, all of which lead up to the same value of the Planck mass at the end of inflation. The predicted value of the Planck mass is consistent with the observed value within a comfortable range of values of the parameters involved. ",
        "introduction": "The extended inflation action is given by \\cite{extended} \\begin{equation} \\label{a} S =\\int d^{4}x \\,\\sqrt{-g}\\left[\\Phi R - {\\omega\\over\\Phi}(\\partial \\Phi)^{2} - \\frac{1}{2}(\\partial \\sigma)^{2} \\\\ - V(\\sigma)\\right] , \\end{equation} where $R$ is the curvature scalar and the potential is $V(\\sigma) = V_0 e^{-\\kappa\\sigma}$. The coupling $\\omega$ plays a similar r\\^{o}le as that of the coupling functions $B_i(\\Psi)$ of the dilaton field $\\Psi$ in string theory. Based on this analogy, several authors have investigated the so-called {\\it hyperextended inflation} models \\cite{extended2,extended3,mikel2}, where $\\omega$ becomes dependent on the BD field. In this paper however, we will merely examine the $\\omega={\\rm const}$ model. The BoI boundary is given by $V(\\sigma)= M^{4}_{\\rm p}(\\Phi)$ or equivalently, \\begin{equation} \\label{c1} \\Phi = \\frac{V_0^{1/2}}{16\\pi}\\,e^{-\\kappa\\sigma/2}. \\end{equation} The BoI is the quantum limit where the metric fluctuations become significant and the inflaton field cannot take the values for which the potential is above this boundary. The EoI boundary is on the other hand \\begin{equation} \\label{d} \\frac{1}{2}\\dot \\sigma^{2} + \\omega\\frac{\\dot \\Phi^{2}}{\\Phi} \\approx V(\\sigma). \\end{equation} The equations of motion in a flat FRW background are \\begin{equation} \\label{e1} \\left(D^2+\\frac{1}{2\\omega}R\\right)\\,\\Phi=0\\,, \\end{equation} \\begin{equation} \\label{e2} D^{2}\\sigma = -V^{\\prime}(\\sigma)\\,, \\end{equation} \\begin{equation} \\label{e3} H^{2}+H\\frac{\\dot \\Phi}{\\Phi} = \\frac{\\omega}{6} \\left({\\dot\\Phi\\over\\Phi}\\right)^2 + \\frac{1}{6\\Phi} \\left[\\frac{1}{2}\\dot\\sigma^{2} + V(\\sigma)\\right] , \\end{equation} and the differential operator $D$ is defined \\begin{equation} \\label{f} D^{2} \\equiv \\partial^{2}_{t} + 3H\\partial_{t} . \\end{equation} In the slow-roll approximation, $\\ddot \\Phi \\ll H\\dot\\Phi \\ll H^{2}\\Phi$ and $\\dot \\sigma^{2}+2\\omega\\, \\dot\\Phi^{2}/\\Phi \\ll 2V(\\sigma)$, (\\ref{e1})-(\\ref{e3}) read \\begin{eqnarray} \\label{g1} \\frac{\\dot\\Phi}{\\Phi} &=&2\\frac{H}{\\omega} \\,, \\\\ \\dot \\sigma &=& -\\frac{1}{3H}V^{\\prime}(\\sigma) \\,,\\\\ H^{2} &=&\\frac{1}{6\\Phi}V , \\end{eqnarray} and the curvature scalar is given by $R=-12H^{2}$. The slow-roll equations (\\ref{g1}) enable us to rewrite (\\ref{d}): \\begin{equation} \\label{h} \\Phi_* =\\left(3-\\frac{2}{\\omega}\\right)\\frac{1}{\\kappa^{2}}, \\end{equation} where the $*$ subindex denotes the value at the end of inflation. Hence, $\\Phi_*$ is independent of $\\sigma$ and, for reasonably large $\\omega$, it is solely determined by the slope of the potential. The condition $\\Phi>0$ also imposes the constraint $\\omega$, as can be seen in Fig.~1, such that the range $0<\\omega<2/3$ is excluded to prevent imaginary values of the Planck mass. The classical trajectories of the fields are given by the following conservation law \\cite{mikel2} \\begin{equation} \\label{h1} \\frac{d}{dt}\\left[\\omega\\Phi + \\int d\\sigma \\frac{V(\\sigma)}{V^{\\prime}(\\sigma)} \\right]=0, \\end{equation} which in the case of the exponential potential yields \\begin{equation} \\label{h2} \\Phi= \\frac{\\sigma}{\\kappa \\omega} + \\left(\\Phi_0- \\frac{\\sigma_0}{\\kappa \\omega}\\right). \\end{equation}  In Fig.~2 we have put together the BoI and EoI curves and the classical trajectories of the fields on the ($\\sigma$,$\\Phi$) plane, i.e. (\\ref{c1}),(\\ref{h}) and (\\ref{h2}) respectively. It can be seen in the figure that inflation takes place in the region enclosed by BoI and EoI to the right of the intersection point $A$. The trajectories given by (\\ref{h2}) are straight lines parallel to the segment $BC$, and the fields drift along these curves in the direction $B\\to C$ during the course of inflation. The region enclosed by BoI and EoI to the left of $A$ does not undergo inflation, because the orientation of the classical trajectories is such that the fields would move from EoI towards BoI, which is not an acceptable solution. In addition to the classical trajectories quantum diffusion is responsible for the jumps of the fields between neighbouring classical trajectories. It can be seen that, unlike with powerlaw potentials, for which $\\sigma$ decreases as $\\Phi$ increases during the course of inflation, in the case of the exponential potential both fields increase during the slow roll.  \\begin{figure}[t] \\centering \\leavevmode\\epsfysize=5.5cm \\epsfbox{omega.ps}\\\\ \\vskip 0.2cm \\caption[fig1]{BD field at EoI, $\\Phi_*$, vs. $\\omega$ for an arbitrary value of $\\kappa$. $\\Phi_*$ is given in units of $\\kappa^{-2}$. Inflation takes place in the range $\\omega<0$ or $\\omega\\geq 2/3$.} \\end{figure}   \\begin{figure}[t] \\centering \\leavevmode\\epsfysize=5.5cm \\epsfbox{fields.ps}\\\\ \\vskip 0.2cm \\caption[fig2]{Predicted BoI and EoI curves, dashed and thick solid curves respectively. Classical trajectories are straight lines parallel to $BC$. At the intersection point $A$ the onset and end of inflation coincide. $\\sigma_{\\rm max}$ determines the scale of validity of the slow-roll approximation. Inflation takes place within the region enclosed by BoI, EoI and $\\sigma\\approx\\sigma_{\\rm max}$.} \\end{figure}   The EoI boundary (\\ref{h}) gives a definite and unique prediction for $\\Phi_*$, and also it implies that if $\\Phi_0>\\Phi_*$ inflation will not occur. In the case of $0<\\Phi_0\\lsim\\Phi_*$, inflation takes place for values of the inflaton \\begin{equation} \\label{sigma0} \\sigma_0 \\gsim -\\frac{2}{\\kappa}\\, {\\rm log}\\left({16\\pi\\Phi_0 \\over V_0^{1/2}}\\right). \\end{equation} Naturally if $\\Phi_0=\\Phi_*$, then the RHS of (\\ref{sigma0}) is $\\sigma_A$, the value of the field at the intersection point $A$ of BoI and EoI in Fig.~2.  It must be noted that the slow-roll approximation does not hold in the case of the exponential potential for arbitrarily large values of $\\sigma$. For a given value of $\\kappa$ it is straightforward to compute $\\sigma_{\\rm max}$ for which the potential and kinetic energy of the fields are comparable and thus the slow-roll conditions break down. It is easy to show from (\\ref{d}) that this scale is \\begin{equation} \\label{smax} \\sigma_{\\rm max}\\approx \\left({3\\omega-2\\over\\kappa}\\right) . \\end{equation} Therefore it follows that the EoI boundary does not span from $\\sigma_A$ to infinity, but inflation will occur within a finite region $\\sigma_A\\lsim\\sigma\\lsim\\sigma_{\\rm max}$. ",
        "conclusions": "In this paper we have examined extended inflation with an exponential potential. The remarkable feature of this model is the prediction of a constant distribution of the Planck mass at the end of inflation, given by (\\ref{h}). The parameter $\\kappa$ of the theory is therefore estimated via the observed Planck mass in this region of the universe, which in turn fixes the parameter $\\sigma_{\\rm max}$ that determines the range of values of $\\sigma$ for which inflation takes place.  The amplitude of the potential $V_0$ is left unconstrained by astrophysical bounds on the spectrum of fluctuations, as described by the argument given in \\S 4. The dynamics as is given in \\S 2 and the likelihood distributions in \\S 3 are shown to be insensitive to the numerical value of this parameter.  As is shown in Fig.~2, the BoI and EoI curves in this model cross at $\\sigma=\\sigma_A$ and the area enclosed between them is thus infinite. However the breakdown of the slow-roll approximation for the exponential potential over the range $\\sigma\\gsim \\sigma_{\\rm max}$ (where $\\sigma_{\\rm max}$ is given by (\\ref{smax})) implies that in practical terms only a finite region of the ($\\sigma$,$\\Phi$) plane undergoes inflation. In the classification of \\cite{V} this means that the exponential potential is {\\it class I}, i.e. the values of the fields at EoI remain finite.      \\subsection*"
    },
    "9804/astro-ph9804239_arXiv.txt": {
        "abstract": "This study is the first known attempt to search for gamma-ray burst repeaters combining data from gamma-ray experiments flying on board different satellites and making use of information derived from the bursts detected simultaneously by all the experiments.  The proposed method is suitable to correlate GRB data provided by experiments that overlap partially or totally in time.  As an application of this method we have correlated the positions of 57 gamma-ray bursts observed by WATCH/GRA\\-NAT and WATCH/EURECA with 1905 bursts detected by BATSE.  Comparing the so-called~~~``added correlation''~~~between~~~the WATCH and BATSE bursts with that obtained with simulated WATCH catalogues, we conclude that there is no indication of recurrent activity of WATCH bursts in the BATSE sample. We derive an upper limit of $15.8\\%$, with a confidence level of $94\\%$, for the number of WATCH gamma-ray bursts that could represent a population of repeaters in the BATSE sample. ",
        "introduction": "Despite the advances carried out so far, the origin of the gamma-ray bursts (hereafter GRBs) remains unknown. The identification of absorption lines in the optical spectrum of GRB 970508 strongly supports models arising from sources at cosmological distances (Metzger et al.  1997), but there is still a lack of knowledge on the mechanisms originating these enigmatic phenomena.  One of the most important clues that could clarify the nature of the GRBs would be the detection of a repeater behaviour.  Initial studies showed an apparent evidence of repetition for the BATSE 1B catalogue (Quashnock and Lamb 1993), suggesting that it would be possible to have an excess of pairs of GRBs clustered in both time and space (Wang and Lingenfelter 1995).  This fact was not confirmed by the work carried out using the BATSE 2B catalogue (Brainerd et al. 1995), although other studies provided marginal evidence for both temporal and angular clustering (Petrosian and Efron 1995). Analyses based on autocorrelations with data from the BATSE 3B catalogue did not find any evidence of repetition (Bennett and Rhie 1996) and have imposed several constraints to the number of repeaters (Tegmark et al.  1996).  Finally, recent studies confirm the lack of repetition in the 4B catalogue and lead to an upper limit to the repetition rate of $ 0.04$ burst source$^{-1}$ yr$^{-1}$ (Hakkila et al. 1997).   The BATSE~4B catalogue was obtained by the BATSE experiment on board the {\\it CGRO} satellite and contains 1637 GRBs detected from April 1991 to August 1996 (Paciesas et al. 1998). The BATSE experiment consists of eight identical detector modules, placed at the corners of the {\\it CGRO} spacecraft and covering energy channels from $\\sim 25$ keV to $\\sim$ 2 MeV. It provides error boxes with a minimum radius of $1.6^{\\circ}$ (1$\\sigma$ confidence level, Fishman et al. 1994). BATSE is detecting bursts at a rate of 0.8 bursts per day. The bursts are daily added to the so-called Current GRB Catalogue, which contains the BATSE~4B catalogue plus all bursts detected after August 1996. When this study was started, the catalogue contained 1905 sources; this sample constitutes the basis of the present work.   The WATCH X-ray all-sky monitor is based on the rotation modulation principle (Lund 1986). The instrument has a circular field of view of 4 steradians and an effective area of $\\sim$ 30 cm$^2$ (averaged over the field of view). Position sensitivity is achieved using the rotation collimator principle, with the collimator grids rotating with a frequency $\\omega$=1 Hz. The phoswich detectors consist of interleaved scintillator-strips of NaI and CsI crystals.  The geometric area of the scintillator is 95~cm$^2$.  Four units were mounted on board the Soviet {\\it GRANAT} satellite in a tetrahedral configuration covering the whole sky, and one unit on board the European Space Agency {\\it EURECA} spacecraft. The total energy range is 8-80~keV, therefore overlapping with the lower BATSE energy band.  WATCH\\-/GRANAT detected bursts in 1990-94 and WATCH/EURECA in 1992-93, thus both experiments also overlapped in time with BATSE. One of the main advantages of WATCH was the capability of locating bursts with relatively small error boxes ($3\\sigma$ error radii with $\\sim$ 1$^{\\circ}$) (Brandt et al. 1990).  WATCH\\-/GRANAT detected 47 GRBs in this period and WATCH\\-/EURECA 12 (Castro-Tirado et al.  1994, Brandt et al. 1994, Sazonov et al. 1998).  Two GRBs (GRB 920814 and GRB 921022) were detected by both the WATCH\\-/GRANAT and WATCH\\-/EURECA experiments. Therefore, the sample of WATCH GRBs used in this study comprises 57 GRBs: 45 WATCH\\-/GRANAT bursts, 10 WATCH\\-/EURECA bursts and the above-mentioned two GRBs.  BATSE also detected 27 of them.  Fig.~\\ref{figure1} shows the sample of 57 WATCH GRBs used in this study. \\begin{figure*} \\centering \\resizebox{!}{!}{\\includegraphics[width=\\hsize,totalheight=10.7cm]{7502.f1}} \\caption{Error boxes for the 57 GRBs detected by WATCH, represented in galactic coordinates. The sample contains 45 GRBs detected by WATCH/GRANAT, 10 by WATCH/EURECA and two localized by both experiments at the same time. The typical radii of the error boxes are $\\sim 1 ^{\\circ}$, with a $3 \\sigma$ confidence level.} \\label{figure1} \\end{figure*}  The distribution of time amplitudes for GRBs shows two classes of bursts: a) durations shorter than $\\sim$ 2 s and b) longer than $\\sim$ 2 s (Kouveliotou et al.  1993). It was noticed that the energy spectra of the short bursts were generally harder than those of the long ones (Kouveliotou et al. 1993, Lestrade et al.  1993).  The fraction of short events in the WATCH sample is smaller than that in the 4B catalogue. This fact can be justified by at least three selection effects:  i) The availability of WATCH for localizing sources is governed by the rotation speed of the collimator grids (1~Hz). So, a source needs to be bright enough for at least one rotation of the modulation collimator in order to be localized, implying a burst duration longer than 1~s. In contrast, the BATSE experiment is able to detect bursts with durations as short as 64 ms.  ii) The low energy band of the WATCH experiment ($\\sim$8-20~KeV) is sensitive to the soft GRBs, below the BATSE lower limit ($\\sim$25 KeV), which generally belong to the class of bursts with durations longer than 2 s.  iii) On the other hand, since WATCH is about an order of magnitude less sensitive than the large-area detectors of BATSE, the WATCH catalogue contains bursts which are brighter than those in the BATSE sample.  The above three reasons~~explain~~why~~the GRBs~~in the WATCH sample are longer, softer and brighter than the average BATSE 4B bursts.  This study is the first known attempt to search for repeaters combining data of $\\gamma$-ray experiments flying on board different satellites. The method proposed makes use of the so-called ``simultaneous bursts'' and is suitable to correlate GRB data provided by experiments that overlap partially or totally in time. In the future, this work could also be used to detect systematic pointing errors between different $\\gamma$-ray experiments, allowing to improve their capability for locating GRBs. ",
        "conclusions": "In this study we have developed a method that allows us to search for GRBs common to two catalogues of sources, each one based on a different instrument.  The method makes use of the GRBs detected simultaneously by both experiments, so it is necessary that the experiments overlap in time. We have applied~~the~~method~~to the WATCH (WATCH/GRANAT + WATCH\\-/EU\\-RE\\-CA) and BATSE (BATSE 4B + bursts detected after August 1996) catalogues.  We conclude that there is no evidence of recurrent activity of WATCH bursts in the BATSE sample. We claim (with a $94\\%$ confidence level) that no more than a $15.8\\%$ of the 57 GRBs detected by WATCH are present in the sample of 1905 BATSE bursts (excluding the simultaneous bursts).  However, the possibility of finding repeaters in each single catalogue cannot be ruled out. Our results support models which do not predict repetitions of GRBs (for instance the merging of neutron stars at cosmological distances)."
    },
    "9804/astro-ph9804180_arXiv.txt": {
        "abstract": "s{ We address the problem of map-making with data from the \\ps\\ High Frequency Instrument, with an emphasis on the understanding and modelling of instrumental effects, and in particular that of sidelobe straylight. } ",
        "introduction": " ",
        "conclusions": "We have shown that the scan strategy of \\ps\\ along rings on the sky allows to decompose the problem of converting data streams into CMB anisotropy maps in two independent steps. This makes the problem tractable numerically, and helps in analysing and monitoring the impact of systematic effects. In particular, a promising method for the identification and removal of sidelobe signals has been developped."
    },
    "9804/astro-ph9804149_arXiv.txt": {
        "abstract": "We review recent results of Sunyaev-Zel'dovich effect (SZE) observations toward galaxy clusters. Using cm-wave receivers mounted on the OVRO and BIMA mm-wave arrays we have obtained high signal to noise images of the effect for more than 20 clusters.  We present current estimates of the Hubble constant and cosmological parameters and discuss the potential of conducting statistical studies with large SZE cluster samples. ",
        "introduction": "Over the last few years there has been a tremendous increase in the study of galaxy clusters as cosmological probes, initially through the use of X-ray emission observations, and in recent years, through the use of Sunyaev-Zel'dovich effect (SZE). Briefly, the SZE is a distortion of the cosmic microwave background (CMB) radiation by inverse-Compton scattering of thermal electrons within the hot intercluster medium (Sunyaev \\& Zel'dovich 1980, see Birkinshaw 1998 for a recent review). The change in the CMB brightness temperature observed is: \\begin{equation} \\frac{\\Delta T}{T_{\\rm CMB}} = \\left[ \\frac{x (e^{x}+1)}{e^{x}-1} -4 \\right] \\int \\left(\\frac{k_B T_e}{m_e c^2}\\right) n_e \\sigma_T dl, \\end{equation} where $x = h \\nu/k_B T_{\\rm CMB}$, and $n_e$, $T_e$ and $\\sigma_T$ are the electron density, electron temperature and the cross section for Thomson scattering. The integral is performed along the line of sight through the cluster.  The other important observable of the hot intercluster gas is the thermal Bremsstrahlung X-ray emission, whose surface brightness $S_X$ can be written as: \\begin{equation} S_X = \\frac{1}{4 \\pi (1+z)^3} \\int n ^{2}_{e} \\Lambda_e dl, \\end{equation} where $z$ is the redshift and $\\Lambda_e(\\Delta E,T_e)$ is the X-ray spectral emissivity of the cluster gas due to thermal Bremsstrahlung within a certain energy band $\\Delta E$. By combining the intensity of the SZE and the X-ray emission observations, and knowing the cluster gas temperature $T_e$, the angular diameter distance to the cluster can be derived due to the different dependence of the X-ray emission and SZE on the electron density, $n_e$. Combining such distance measurements with redshift allows a determination of the Hubble constant, H$_0$, as a function of certain cosmological parameters (e.g., Hughes \\& Birkinshaw 1998a). If distance measurements for a sample of clusters exist, then the angular diameter distance with redshift relation can be used to put constraints on the cosmological models, similar to current supernovae constraints at high redshift. ",
        "conclusions": ""
    },
    "9804/astro-ph9804133_arXiv.txt": {
        "abstract": "A new method of shower-image analysis is presented which appears very powerful as applied to those Cherenkov Imaging Telescopes with very high definition imaging capability. It provides hadron rejection on the basis of a single cut on the image shape, and simultaneously determines the energy of  the electromagnetic shower and the position of the shower axis with respect to the  detector. The source location is also reconstructed for each individual $\\gamma$-ray shower, even with one single telescope, so for a point source the hadron rejection can be further improved. As an example, this new method is  applied to data from the C{\\small AT} (Cherenkov Array at Th\\'emis) imaging  telescope, which has been operational since Autumn, 1996. ",
        "introduction": "Many of the sources in the E{\\small GRET} catalogue \\cite{egret} have well-identified radio, optical, or X-ray counterparts, which give the source position to an accuracy much better than can be achieved by {\\small ACT}  telescopes.  Such sources with known position are usually placed  at the centre of the field of an imaging telescope.  However, many unidentified sources in the  E{\\small GRET} catalogue are located in error boxes with typical size of 1$^\\circ$.  Ground-based  Cherenkov detectors are, in principle, able to localize such sources  with a higher precision.  In order to observe these sources,  different methods have been developed by imaging Cherenkov telescope  groups. Stereoscopy is the most direct way to find the direction  of a source, but requires at least two telescopes and reduces  the collection area \\cite{hegratel}.  For single-telescope experiments with sufficient background rejection on an image-shape criterion, it is possible  to perform de-localized analyses assuming the source at the different points of a grid covering the field of view or by examining the distribution of the intersections  of the image axes \\cite{grid}. However, an event-by-event analysis method  such as that described here is  preferable since in the former methods the signal is more easily drowned-out by the background.  The method has been tested on simulated  $\\gamma$-images provided by the Monte-Carlo simulation  program described above and a realistic simulation of the detector response, including the measured variation in collection efficiency and gain  between the phototubes, and measured wavelength response of the mirrors and Winston cones. For optimization of the cut values, these simulated $\\gamma$-images have been used together with the real background  events from data from off-source runs with the C{\\small AT} imaging telescope. Gammas from a point-like source with the Crab nebula spectrum \\cite{catcrab} (see equation (1) above) were simulated at various elevations.  The capability of the method both  for source detection and for source spectrum measurement have been examined.  \\subsection{Source detection}  As applied to the data, the method consists of minimizing the  $\\chi^2$ with respect to $E_\\gamma$, $D$, $\\vec\\xi$, and $\\phi$. Fig~\\ref{fig:chi2} shows the $\\chi^2$ probability distributions obtained  with this fit for the simulated $\\gamma$-ray events and real background events. \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[0 25 590 520]{fig6.eps} \\caption{$\\chi^2$ probability distribution for a fit with the source coordinates considered as free parameters (constraint is from shape alone): a) vertical $\\gamma$-ray showers; b) real background (off-source data) showers. The upper line in each figure is for all events above threshold, the shaded distributions for events with a fitted energy, $E_{\\mathrm{f}}$, greater than $350\\:{\\mathrm {GeV}}$ and a fitted impact parameter, $D_{\\mathrm{f}}$, between 30 and $125\\:{\\mathrm{m}}$.} \\label{fig:chi2} \\end{figure} A cut on the $\\chi^2$ probability value, $P(\\chi^2)$, provides a selection of $\\gamma$-like events on the basis of the image shape alone. The reconstructed angular  origins obtained for simulated $\\gamma$-events accumulate around the actual source position which in this case is at the centre of the field.  The dispersion around the actual source position  has a typical {\\small RMS} spread of $0.14^\\circ$.  For each event, the accuracy of the angular origin determination is better by a factor two in the direction perpendicular to the image axis than in the direction of the image axis (Fig.~\\ref{fig:pointerr}). \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[20 25 540 230]{fig9.eps} \\caption{Distributions of longitudinal and transverse errors (with respect to the image major axis) in the reconstruction of the source position for vertical showers with $P(\\chi^2)>0.2$.  The small bump at negative values of longitudinal error results from wrong direction reconstruction (mainly for events close to the energy threshold). Shaded histograms correspond to events with a fitted energy greater than $350\\:{\\mathrm GeV}$.} \\label{fig:pointerr} \\end{figure} The {\\small RMS} longitudinal error typically varies  from $0.2^\\circ$ to $0.1^\\circ$ as the energy varies from the threshold to $2\\:{\\mathrm {TeV}}$.  The angular origins obtained for background events are spread over the whole field with an approximately Gaussian distribution with a $1.8^\\circ$ {\\small FWHM}. Since this distribution is fairly flat, 2-dimensional skymaps of the  angular origins of the showers could be used for source detection, as can be seen from the reconstructed positions in data taken on Markarian 501 in Fig.~\\ref{fig:m501}. \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[0 25 580 520]{fig10.eps} \\caption{The distribution of reconstructed angular origins for the data from a 30-minute run on Markarian 501 from April 16, 1997, for events with $P(\\chi^2)>0.2$. The concentric circles represent the trigger region and the small-pixel region, respectively.  During the run, the optic axis described the small arc indicated due to the mechanical flexibility of the structure, which is monitored as  described in [9]. The number of events reconstructed in each bin of $(0.05{^0})^2$ is shown.  No  background subtraction has been performed.} \\label{fig:m501} \\end{figure} The errors in angular reconstruction given above are for a single shower; a point source with poorly defined position could thus be localized to $\\sim 1-2'$ with the combination of $\\sim 100$ such events.  For the present, a conservative procedure of monitoring the background is used,  based on the pointing angle $\\alpha$, similar to the ``orientation'' of Whipple  \\cite{scuts}: $\\alpha$ is the angle at the image barycentre between the actual  source position and the reconstructed source position. The pointing angle does  not use the full information contained in the results of the fit, but has a  fairly flat distribution from $0^\\circ$ to about $120^\\circ$ for background events, which allows the background level to be easily monitored.   The cut on $\\alpha$ is more efficient than a cut on the angular distance between the source position and the reconstructed $\\gamma$ origin since, as seen in  Fig.~\\ref{fig:pointerr}, the position reconstructed is not symmetric about the source position. The distribution of $\\alpha$ for $\\gamma$-events from a Crab-like source exhibits a peak at $0^\\circ$ (Fig.~\\ref{fig:alpha}) and a small accumulation at $180^\\circ$ corresponding to events which are wrongly found to point away from the source. \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[5 25 580 520]{fig7.eps} \\caption{Distribution of the pointing angle $\\alpha$ for events with $P(\\chi^2)>0.2$ (constraint from shape alone): a) vertical $\\gamma$-ray showers; b) Real background (off-source data) showers. The upper line is for all events above threshold, the shaded histograms for events with a fitted energy greater than $350\\:{\\mathrm {GeV}}$.} \\label{fig:alpha} \\end{figure} Around 17\\% of the $\\gamma$-events from a Crab-like source are in this situation.  The proportion of events with a wrongly reconstructed direction decreases with increasing energy, from 22\\% at $200\\:{\\mathrm {GeV}}$ to 9\\% at $600\\:{\\mathrm {GeV}}$ and 4\\% at $1\\:{\\mathrm {TeV}}$.  The significance of a signal is calculated using the usual formula: $(ON-OFF)/\\sqrt{ON+OFF}$ \\cite{lima}, assuming equal time on and off-source.  The significance per hour on a simulated Crab-like source at zenith has been calculated for various cut values on $\\alpha$ and $P(\\chi^2)$ (Fig.~\\ref{fig:signif1}.a). \\begin{figure}[t] $$ \\epsfxsize=6.5 cm (a)\\epsffile[40 30 540 490]{fig8.eps} \\; \\epsfxsize=6.5 cm (b)\\epsffile[40 30 540 490]{fig11.eps} $$ \\vspace{-1cm} \\caption{Detection significance for a Crab-like source at zenith in one hour of observation, shown as a function of the two cuts on $P(\\chi^2)$ and $\\alpha$: (a) for a source at the centre of the camera; (b) for a source at $1^\\circ$ from the centre. The full lines indicate contours of equal significance; dotted lines show fixed $\\gamma$-ray selection efficiency.} \\label{fig:signif1} \\end{figure} The best result in terms of both significance and efficiency for $\\gamma$-events is obtained for a $P(\\chi^2)>0.2$ and $\\alpha<6^\\circ$, which gives $5.1\\sqrt{t}\\ \\sigma$ where $t$ is the on-source observation time in hours, retaining 34\\% of the $\\gamma$-events  while giving a rejection factor of 120 on background events.  This rejection factor is smaller than for some comparable  experiments as there is a large rejection factor at the trigger level, allowing a moderate background rate of $15\\:{\\mathrm {Hz}}$ at the zenith. At $45^\\circ$ from zenith the best significance for the same cuts falls to $2.7\\sqrt{t}\\ \\sigma$. This is essentially due to the higher energy threshold, leading to a lower event rate; on the other hand, for the same $\\gamma$-ray selection efficiency the background rejection factor is comparable to that at zenith.  In order to estimate the efficiency of the $\\chi^2$-method in the case of a source with a poorly-defined position, a simulated Crab-like source has been set on the edge of the trigger area ($1^\\circ$ from the centre). In this case, the equivalent detection area is divided by a factor of the order of two.  Even for a source with known position, when the source is not at the centre of the field it is possible to use one side of the camera as the off region for the other side \\cite{offc}.  For a Crab-like source at zenith, the same cuts in $P(\\chi^2)$ and $\\alpha$ as for a source in the centre of the camera give a $2.3\\sqrt{t}\\ \\sigma$ significance, with a selection efficiency for $\\gamma$'s of 36\\% and a rejection factor of 116 above threshold (Fig.~\\ref{fig:signif1}.b). This means that the on-source run time has to be five times larger for a source on the edge of the trigger area than for a source at the centre of the field for the same significance.  The corresponding significance obtained at $45^\\circ$ from zenith for an  off-centre source is $1.5\\sqrt{t}\\ \\sigma$.  \\subsection{$\\gamma$-ray energy measurement}  For a source detected with  a strong enough significance, the energy spectrum can be studied by a detector with good energy resolution.  The fit described above in which the source position is a free parameter gives a first estimate of the energy of each event to within about 30\\%. However, more precise spectral studies can be carried out on point sources of $\\gamma$-rays. The use of the source position as a constraint in the fit provides a higher accuracy for impact parameter measurement and, as a consequence, for energy measurement. If trigger selection effects are ignored in the Monte-Carlo program (thus accepting all events above $100\\:{\\mathrm {GeV}}$), the $\\chi^2$ minimization with respect to $E_\\gamma$, $D$, and $\\phi$  provides an unbiased energy measurement within about 25\\% (statistical error only). Trigger selection effects are small for events well above the threshold, as can be seen for simulated $400\\:{\\mathrm {GeV}}$ $\\gamma$-rays in Fig.~\\ref{fig:ener500}. \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[15 25 550 300]{fig12.eps} \\caption{Distribution of $\\ln({E_{\\mathrm {f}}}/{E_\\gamma})$ for vertical $400\\:{\\mathrm {GeV}}$ $\\gamma$-ray showers satisfying the selection cuts (the source location, $\\vec{\\xi}$, being fixed in the fit). The shaded histogram is further restricted to events with a fitted impact parameter $30\\:{\\mathrm {m}} < D_{\\mathrm {f}} < 125\\:{\\mathrm {m}}$.} \\label{fig:ener500} \\end{figure} This figure also shows that the distribution of the fitted event energies about the true energy is Gaussian on a logarithmic scale. Consequently, the slope of a power-law spectrum can be directly estimated with this technique. Close to the threshold, however, the fitted energy $E_{\\mathrm {f}}$ is overestimated as a consequence of the trigger selection. Similarly, the small remaining bias in $ \\log (E_{\\mathrm {f}}/E_\\gamma)$ at $400\\:{\\mathrm {GeV}}$ is due to showers with large impact parameters for which the trigger selection is critical at this energy since the telescope is then located at the border of the light pool (Fig.~\\ref{fig:densite}). This effect is largely removed if only showers with a fitted  impact parameter $D_{\\mathrm {f}}$ lower than $125\\:{\\mathrm {m}}$ are included (shaded histogram in Fig.~\\ref{fig:ener500}). The bias induced by the trigger selection at different energies is best illustrated  by plotting 68\\% confidence intervals for $E_{\\mathrm {f}}$ as a function of the true value $E_\\gamma$ used in the simulation (Fig.~\\ref{fig:enerfit}). \\begin{figure} \\epsfxsize=14.5 cm \\leavevmode \\centering \\epsffile[5 25 580 520]{fig14.eps} \\caption{Fitted energy, $E_{\\mathrm {f}}$, versus true energy, $E_\\gamma$, for $\\gamma$-ray showers satisfying the selection cuts (the source location, $\\vec{\\xi}$, being fixed in the fit) and for which $30\\:{\\mathrm {m}}< D_{\\mathrm {f}}\\cos Z < 125\\:{\\mathrm {m}}$ ($Z=$~zenith angle). The shaded interval shows the 68\\%  confidence intervals for a source at the zenith.  The effect of trigger selection on the energy estimation can be seen.} \\label{fig:enerfit} \\end{figure} It can be seen that for $E_f$ below $350\\:{\\mathrm {GeV}}$ for vertical showers, only an upper limit can be safely derived for $E_\\gamma$. Therefore, spectrum measurement is reliable only above a spectrometric threshold, that is, in the region in which $E_{\\mathrm {f}}$ depends linearly on $E$. It is somewhat higher than the nominal threshold which is relevant for source discovery. Fig.~\\ref{fig:enerfit} also shows the variation of the average  values of $\\log(E_{\\mathrm {f}})$ on $\\log(E_\\gamma)$ for zenith angles $0^{\\circ}$,  $30^{\\circ}$, $45^{\\circ}$, and $60^{\\circ}$, showing the increase in the spectrometric thresholds with increasing zenith angle. Restricting to events with $E_{\\mathrm {f}}$ above the spectrometric threshold and $30\\:{\\mathrm {m}}<  D_{\\mathrm {f}} \\cos Z <  125\\:{\\mathrm {m}}$, the accuracy of the preceding method is about 20\\% (statistical error only), independent of the zenith angle $Z$ up to $45^{\\circ}$. ",
        "conclusions": "The method described above, based on a realistic analytic description of electromagnetic air showers, is best suited for those Cherenkov Imaging Telescopes with a high-resolution camera. The light distribution in the focal plane is fully exploited, yielding the shower direction from the asymmetry of the longitudinal profile as well as the source position in the focal plane. Selection of $\\gamma$-rays on the basis of the image shape is performed by using a single $\\chi^2$-variable instead of a series of cuts on various image parameters.  By combining the $\\chi^2$ probability cut and a direction ($\\alpha$) cut, a significance of 5$\\sigma$ per hour can be achieved for a Crab-like source at the zenith. A future development of the method would be to use the distribution of selected $\\gamma$-ray origins on the celestial sphere together with the known energy-dependent point spread function of the method to estimate the significance by a maximum-likelihood method. With the C{\\small AT} telescope ($250\\:{\\mathrm {GeV}}$ threshold), sources with the intensity of the Crab nebula can be  detected in one hour.  This has been confirmed with the results obtained in the  96/97  observing campaign.  Moreover, sources with poorly defined position can be localized with an accuracy of the order of an arc minute on the basis of about  100 showers. The accuracy on $\\gamma$-ray energy is of the order of 20--25\\%. Biases induced by the trigger selection have been investigated; in particular, care should be taken in spectrum measurement, which is accurate only above a specific threshold, higher than that used for source detection."
    },
    "9804/astro-ph9804305_arXiv.txt": {
        "abstract": "We analyze the stability of g-modes in variable white dwarfs with hydrogen envelopes. All the relevant physical processes take place in the outer layer of hydrogen rich material which consists of a radiative layer overlain by a convective envelope. The radiative layer contributes to mode damping because its opacity decreases upon compression and the amplitude of the Lagrangian pressure perturbation increases outward. The convective envelope is the seat of mode excitation because it acts as an insulating blanket with respect to the perturbed flux that enters it from below.  A crucial point is that the convective motions respond to the instantaneous pulsational state. Driving exceeds damping by as much as a factor of two provided $\\omega\\tau_c\\geq 1$, where $\\omega$ is the radian frequency of the mode and $\\tau_c\\approx 4\\tau_{\\rm th}$ with $\\tau_{\\rm th}$ being the thermal time constant evaluated at the base of the convective envelope. As a white dwarf cools, its convection zone deepens, and modes of lower frequency become overstable.  However, the deeper convection zone impedes the passage of flux perturbations from the base of the convection zone to the photosphere. Thus the photometric variation of a mode with constant velocity amplitude decreases. These factors account for the observed trend that longer period modes are found in cooler DAVs. Overstable modes have growth rates of order $\\gamma\\sim 1/(n\\tau_\\omega)$, where $n$ is the mode's radial order and $\\tau_\\omega$ is the thermal time-scale evaluated at the top of the mode's cavity. The growth time, $\\gamma^{-1}$, ranges from hours for the longest period observed modes ($P\\approx 20$ minutes) to thousands of years for those of shortest period ($P\\approx 2 $ minutes).  The linear growth time probably sets the time-scale for variations of mode amplitude and phase.  This is consistent with observations showing that longer period modes are more variable than shorter period ones.  Our investigation confirms many results obtained by Brickhill in his pioneering studies of ZZ Cetis. However, it suffers from at least two serious shortcomings. It is based on the quasiadiabatic approximation that strictly applies only in the limit $\\omega\\tau_c\\gg 1$, and it ignores damping associated with turbulent viscosity in the convection zone. We will remove these shortcomings in future papers. ",
        "introduction": "}  ZZ Cetis, also called DAVs, are variable white dwarfs with hydrogen atmospheres.  Their photometric variations are associated with nonradial gravity-modes (g-modes); for the first conclusive proof, see Robinson \\etal (\\cite{scaling-robinson82}). These stars have shallow surface convection zones overlying stably stratified interiors. As the result of gravitational settling, different elements are well separated . With increasing depth, the composition changes from hydrogen to helium, and then in most cases to a mixture of carbon and oxygen. From center to surface the luminosity is carried first by electron conduction, then by radiative diffusion, and finally by convection.  Our aim is to describe the mechanism responsible for the overstability of g-modes in ZZ Ceti stars. This topic has received attention in the past. Initial calculations of overstable modes were presented in Dziembowski \\& Koester (\\cite{adia-dziem81}), Dolez \\& Vauclair (\\cite{adia-dolez81}), and Winget \\etal (\\cite{adia-winget82}). These were based on the assumption that the convective flux does not respond to pulsation; this is often referred to as the frozen convection hypothesis.  Because hydrogen is partially ionized in the surface layers of ZZ Ceti stars, these workers attributed mode excitation to the $\\kappa$-mechanism. In so doing, they ignored the fact that the thermal time-scale in the layer of partial ionization is many orders of magnitude smaller than the periods of the overstable modes. Pesnell (\\cite{adia-pesnell87}) pointed out that in calculations such as those just referred to, mode excitation results from the outward decay of the perturbed radiative flux at the bottom of the convective envelope. He coined the term `convective blocking' for this excitation mechanism.\\footnote{This mechanism was described in a general way by Cox \\& Guili (\\cite{adia-cox68}), and explained in more detail by Goldreich \\& Keeley (\\cite{adia-goldreich77}).} Although convective blocking is responsible for mode excitation in the above cited references, it does not occur in the convective envelopes of ZZ Ceti stars. This is because the dynamic time-scale for convective readjustment (i.e., convective turn-over time) in these stars is much shorter than the g-mode periods. Noting this, Brickhill (\\cite{adia-brick83}, \\cite{adia-brick90}, \\cite{adia-brick91a}, \\cite{adia-brick91b}) assumed that convection responds instantaneously to the pulsational state. He demonstrated that this leads to a new type of mode excitation, which he referred to as convective driving. Brickhill went on to presents the first physically consistent calculations of mode overstability, mode visibility, and instability strip width. Our investigation supports most of his conclusions. Additional support for convective driving is provided by Gautschy \\etal (\\cite{adia-gautschy96}) who found overstable modes in calculations in which convection is modeled by hydrodynamic simulation.  In this paper we elucidate the manner in which instantaneous convective adjustment promotes mode overstability. We adopt the quasiadiabatic approximation in the radiative interior. We also ignore the effects of turbulent viscosity in the convection zone. These simplifications enable us to keep our investigation analytical, although we appeal to numerically computed stellar models and eigenfunctions for guidance.  The DA white dwarf models we use are those produced by Bradley (\\cite{scaling-bradley96}) for asteroseismology.  Fully nonadiabatic results, which require numerical computations, will be reported in a subsequent paper. These modify the details, but not the principal conclusions arrived at in the present paper.  The plan of our paper is as follows. The linearized wave equation is derived in \\S \\ref{sec:adia-prepare}. In \\S \\ref{sec:adia-perturb}, we evaluate the perturbations associated with a g-mode in different parts of the star. We devote \\S \\ref{sec:adia-driving} to the derivation of a simple overstability criterion. Relevant time-scales and the validity of the quasiadiabatic approximation are discussed in \\S \\ref{sec:adia-discussion}. The appendix contains derivations of convenient scaling relations for the dispersion relation, the WKB eigenfunction, and the amplitude normalization. ",
        "conclusions": "}   \\subsection{Time-Scales \\label{sec:adia-relevant}}  Three time-scales are relevant for convective driving in DAVs.  The first is the period of an overstable g-mode, $P = 2\\pi/\\omega$, which is typically of order a few hundred seconds. The second is the dynamical time constant, $\\tcv \\sim H_p/\\vcv$, on which convective motions respond to perturbations; $\\tcv \\leq 1 \\s$ throughout the convection zones of even the coolest ZZ Cetis. This is why the convective motions adjust to the instantaneous pulsational state. The third is the thermal time constant, $\\tau_c$, during which the convection zone can bottle up flux perturbations that enter it from below.  Given the central role of $\\tau_c$, we elaborate on its relation both to $\\tcv$ and to the more conventional definition of thermal time constant, $\\tau_{\\rm th}$, at depth $z$. The latter is the heat capacity of the material above that depth divided by the luminosity. In a plane parallel, fully ionized atmosphere this is equivalent to \\begin{equation} \\tau_{\\rm th}\\equiv {1\\over F} \\int_0^z dz \\, c_p\\, {\\rho k_B\\over m_p T} \\approx {5pz\\over 7F}. \\label{eq:adia-tauth}\\end{equation} Appeal to equation \\refnew{eq:adia-Fvcv} establishes that inside the convection zone \\begin{equation} {\\tcv\\over \\tau_{\\rm th}}\\sim \\left({\\vcv\\over c_s}\\right)^2\\ll 1. \\end{equation} Now $\\tau_c\\equiv (B+C)\\tau_b$, where $\\tau_b$ is defined by equation \\refnew{eq:adia-taub}.  To the extent that $c_p\\approx 5$ is constant in the convection zone, $\\tau_b\\approx \\tau_{\\rm th}/5$, where the latter is evaluated at $z_b$.  Next we address the relation between $\\tau_c$ and $\\tau_b$. Here we are concerned with the relatively large value of $B+C$, typically about 20 for DAVs.\\footnote{In our models, $B$ and $C$ have comparable value.} Recall from equations \\refnew{eq:adia-Dsb} and \\refnew{eq:adia-DFph} that \\begin{equation} {\\delta F_{\\ph} \\over F}\\approx {\\delta s_b\\over B+C}. \\label{eq:adia-dFds}\\end{equation} So the photosphere and superadiabatic layer add an insulating blanket on top of the convection zone.  The large value of $B$ follows because the photospheres of DAVs are composed of lightly ionized hydrogen. In this state, the values of $\\kappa_T$ and $s_T$ are both large and positive; typical values in the middle of the instability strip are $\\kappa_T\\approx 6$ and $s_T\\approx 24$.  The large and positive $\\kappa_T$ arises because the population of hydrogen atoms in excited states which the ambient radiation field can photoionize increases exponentially with increasing $T$. The large and positive $s_T$ occurs as the ionization fraction increases exponentially with increasing $T$, and the much larger entropy contributed by a free as compared to a bound electron.  The large and positive value of $C$ reflects the increase in entropy gradient that accompanies an increase in convective flux. It is obtained from mixing length theory with an unperturbed mixing length.  \\subsection{Validity of the Quasiadiabatic Approximation}  The validity of the quasiadiabatic approximation requires that the nonadiabatic parts of the expressions for $\\delta \\rho/\\rho$ and $\\delta T/T$, as given by equations \\refnew{eq:adia-eqdelrho} and \\refnew{eq:adia-eqdelT}, be small in comparison to the adiabatic parts. Thus the ratio \\begin{equation} {\\cal R_{\\rm na}}\\equiv {\\delta s\\over\\delta p/p} \\label{eq:adia-rnonad}\\end{equation} is a quantitative measure of nonadiabaticity. We estimate ${\\cal R}_{\\rm na}$ for the radiative interior and the convection zone.  We calculate $\\delta s$ in the radiative interior from \\begin{equation} \\delta s \\approx {iF\\over \\omega}{m_p\\over\\rho k_B T}{d\\over dz}\\left({\\delta F\\over F}\\right). \\label{eq:adia-delsrad}\\end{equation} In the upper evanescent layer, $z_b<z<z_\\omega$, this leads to \\begin{equation} |{\\cal R_{\\rm na}}|\\sim {1\\over \\omega\\tau_{\\th}}, \\label{eq:adia-ratioevu}\\end{equation} whereas in the propagating cavity, $z>z_\\omega$, we find \\begin{equation} |{\\cal R_{\\rm na}}|\\sim {1\\over \\omega\\tau_{\\th}}\\left({z\\over z_\\omega}\\right). \\label{eq:adia-ratioevl}\\end{equation} Nonadiabatic effects in the radiative interior are maximal at $z=z_b$, where \\begin{equation} |{\\cal R_{\\rm na}}|\\sim {1\\over \\omega\\tau_b}, \\label{eq:adia-ratioevmax}\\end{equation} since $\\tau_{\\rm th}/z$ increases with depth.  The measure of nonadiabaticity in the convection zone is given by equation \\refnew{eq:adia-Dsb}. Since $\\omega\\tau_b>1$ is required for the validity of the quasiadiabatic approximation in the radiative zone, we restrict consideration to the limiting case $\\omega\\tau_c\\gg 1$. In this limit, equation \\refnew{eq:adia-Dsb} yields \\begin{equation} |{\\cal R_{\\rm na}}|\\sim {1\\over \\omega\\tau_b}, \\label{eq:adia-ratiocvz}\\end{equation} which is identical to the value arrived at for the radiative zone in equation \\refnew{eq:adia-ratioevmax}.  The requirement $\\omega\\tau_b\\gtrsim 1$ for the validity of the quasiadiabatic approximation severely limits the applicability of the current investigation. The perturbed flux at the photosphere is related to that at the bottom of the convection zone by \\begin{equation} {\\Delta F_{\\ph}\\over F}\\approx {1\\over 1-i\\omega\\tau_c}{\\Delta F_b\\over F}. \\label{eq:adia-relDelF}\\end{equation} Since $\\tau_c$ is at least an order of magnitude larger than $\\tau_b$, modes with $\\omega\\tau_b\\gtrsim 1$ are likely to exhibit small photometric variations. However, this may not render them undetectable because their horizontal velocity perturbations pass undiminished through the convection zone.  \\subsection{Brickhill's Papers\\label{subsec:adia-brickhill}}  Our investigation is closely related to studies of ZZ Cetis by Brickhill (\\cite{adia-brick83},\\cite{adia-brick90},\\cite{adia-brick91a}). Brickhill recognized that the convective flux must respond to the instantaneous pulsational state. To determine the manner in which the convection zone changes during a pulsational cycle, he compared equilibrium stellar models covering a narrow range of effective temperature. Brickhill provided a physical description of convective driving and obtained an overstability criterion equivalent to ours.\\footnote{Our time constant $\\tau_c$ is equivalent to the quantity $D$ which Brickhill defined in equation (9) of his 1983 paper.} Moreover, he recognized that the convection zone reduces the perturbed flux and delays its phase.  Our excuses for revisiting this topic are that Brickhill's papers are not widely appreciated, that our approach is different from his, and that our paper provides the foundation for future papers which will examine issues beyond those he treated.   \\begin{appendix}"
    },
    "9804/astro-ph9804075_arXiv.txt": {
        "abstract": "The decay of massive neutrinos to final states containing only invisible particles is poorly constrained experimentally. In this letter we describe the constraints that can be put on neutrino mass and lifetime using CMBR measurements. We find that very tight lifetime limits on neutrinos in the mass range 10 eV - 100 keV can be derived using CMBR data from upcoming satellite measurements. \\\\ Keywords: Neutrino decay, Cosmology: Theory, Cosmic Microwave Background \\\\ PACS: 13.35.Hb, 14.60.St, 98.70.Vc ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804243_arXiv.txt": {
        "abstract": "Although the unified formula for $\\gamma$-ray absorption process involving both the magnetic field and a perpendicular electric field derived by Daugherty \\& Lerche (1975) is correct, we argued in this paper that their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading. The key point is that usually the direction of a $\\gamma$ photon at the emission point observed in the laboratory frame should be $(v/c, 0, [1-(v/c)^2]^{1/2})$ rather than $(0, 0, 1)$, where $v$ is the co-rotating velocity. This emission direction is just the one which results in zero attenuation coefficient of the $\\gamma$ photon. Calculation shows that after the photon has moved a distance, its direction lead to the result that the induced electric field is also of minor importance. Thus only $\\gamma-B$ process is the important mechanism for the pair production in the pulsar magnetospheres. The implications of the modification by ejecting the induced electric field are also discussed. ",
        "introduction": "Pair production process plays an important role in pulsar physics. It is not only a necessary process for the multiplication of the particles to account for emissions of different bands from pulsars (e.g. Sturrock 1971; Ruderman \\& Sutherland 1975; Arons \\& Scharleman 1979; Arons 1983; Cheng, Ho, \\& Ruderman 1986), but also an important mechanism to absorb $\\gamma$-rays produced in the pulsar magnetospheres, especially near the polar cap region (e.g. Hardee 1977; Harding, Tademaru, \\& Esposito 1978; Harding 1981; Daugherty \\& Harding 1982, 1996; Zhao {\\it et al.} 1989; Lu \\& Shi 1990; Lu, Wei, \\& Song 1994; Dermer \\& Sturner 1994; Sturner, Dermer, \\& Michel 1995; Wei, Song, \\& Lu 1997). Furthermore, the way by which the $\\gamma$-rays are absorbed is also the key factor to limit the parameters of the inner magnetospheric accelerators of pulsars (e.g. Ruderman \\& Sutherland 1975, hereafter RS75; Zhang \\& Qiao 1996; Qiao \\& Zhang 1996; Zhang {\\it et al.} 1997a; Zhang, Qiao, \\& Han 1997b, hereafter ZQH97b).  Pair formation in intense magnetic fields ($\\gamma-B$ process) has been studied explicitly by different authors (e.g. Erber 1966; Tsai \\& Erber 1974; Daugherty \\& Harding 1983, Rifert, M\\'{e}sz\\'{a}ros \\& Bagoly 1989), and its importance in pulsar physics was first pointed out by Sturrock (1971). Daugherty \\& Lerche (1975, hereafter DL75) first dealt with the case involving a relatively weaker electric field perpendicular to the magnetic field (${\\bf E}^2-{\\bf B}^2\\le0$, ${\\bf E\\cdot B}=0$), and came to a unified formula of the attenuation coefficient of the $\\gamma$ photons. The more general case involving both the perpendicular and the parallel components of the electric field with respect to the magnetic field ($E_\\perp$ and $E_\\parallel$) was presented by Daugherty \\& Lerche (1976) and Urrutia (1978). In the specific case of pulsars, although $E_\\parallel$ is usually sufficiently small so that its effect is negligible, $E_\\perp$ induced by the fast spin of the neutron stars was demonstrated to be very important in pair formation process by DL75. This leads many authors to take this effect seriously into account in their studies (e.g. Hardee 1977; Lu \\& Shi 1990; Lu, Wei, \\& Song 1994; Qiao \\& Zhang 1996).  In this paper, we'll argue that although the unified formula of DL75 is correct, their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading. The detailed argument is presented in Section 2 and Section 3. Finally, we discuss the possible implications of this modification. ",
        "conclusions": "Although the unified formula for the pair production process involving both the magnetic field and the perpendicular electric field derived by DL75 is correct, we have argued in this paper that their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading at least for the ``aligned rotator'' case. At the emission point, the photon emitted by a certain mechanism (e.g. the curvature radiation or the inverse Compton scattering) just moves along the very direction in which the attenuation coefficient is zero. Considering the propagation of the photon, we found that this rotation-induced electric field still plays a minor role in the $\\gamma$-ray absorption process in the polar cap region of a pulsar.  For the general case of an ``oblique rotator'' in which the magnetic and the rotational axes are misaligned, the co-rotating velocity is no more perpendicular to the magnetic field so that $\\theta_B^{\\prime}$ (also $\\theta_u^{\\prime}$ and $\\theta_\\gamma^{\\prime}$) is not $\\pi/2$. The photon direction consequently deviate from $(v/c, 0, [1-(v/c)^2]^{1/2})$ slightly. From Fig.2, we see that the attenuation coefficient is also very small around the direction $(v/c, 0, [1-(v/c)^2]^{1/2})$, so that the conclusion that the induced electric field plays a minor role in $\\gamma$-ray absorption still holds for the oblique rotator case. Actually, DL75's result can only be applied to the aligned case strictly, since generally the co-rotating velocity ${\\bf v}_{r}$ is not equal to ${\\bf v}_{drift}=c({\\bf E\\times B})/B^2$, with which one can define a frame where the electric field vanishes completely. In the co-rotating frame of an oblique rotator, an electric field component parallel to the magnetic field will still remain so that DL75's application condition fails.  Our results in this paper may have some implications for some previous studies which regard the electric field as the important effect of $\\gamma$-ray absorption (e.g. Hardee 1977; Zhao {\\it et al.} 1989; Lu \\& Shi 1990; Lu, Wei, \\& Song 1994; Qiao \\& Zhang 1996).  Although the polar cap models of the $\\gamma$-ray pulsars are by all means sound in principle, their concrete details will alter much by ejecting the induced electric field. Specifically, based on DL75's result, Hardee (1977) got an absolute upper limit to the photon energies $$E_\\gamma <9.6\\times 10^9 B_{12}^{-1}r_6^2 P {\\rm eV}, \\eqno(6)$$ (his Eq.(38)) at which escape of the $\\gamma$-rays from the magnetosphere is possible. It should be replaced by a threshold in the pure magnetic field absorption scheme $$E_\\gamma <2.6\\times 10^9 B_{12}^{-1} P^{1/2} {\\rm eV} \\eqno(7)$$ (Wei, Song, \\& Lu 1997, their Eq.(10), $r_6=1$ and the last open field line is assumed). Thus the generation order parameters proposed by Lu, Wei, \\& Song (1994) should take the form in Wei, Song, \\& Lu (1997, their Eq.(17)).  The three boundary lines (birth line, death line and appearance line) in the $\\dot P-P$ diagram of pulsars derived by Qiao \\& Zhang (1996) are also based on the electric field absorption. The details will also be changed by ejecting the electric field, but the picture still remains and may give a hint to us about the magnetic field configuration in the neutron star vicinity (Qiao \\& Zhang, discussions)."
    },
    "9804/astro-ph9804257_arXiv.txt": {
        "abstract": "Molecules dominate the cooling function of neutral metal-poor gas at high density. Observation of molecules at high redshift is thus an important tool toward understanding the physical conditions prevailing in collapsing gas. Up to now, detections are sparse because of small filling factor and/or sensitivity limitations. However, we are at an exciting time where new capabilities offer the propect of a systematic search either in absorption using the UV Lyman-Werner H$_2$ bands or in emission using the CO emission lines redshifted in the sub-millimeter. ",
        "introduction": "\\subsection{Introduction} QSO absorption line systems probe the baryonic matter over most of the history of the Universe (0~$<$~$z$~$<$~5).  The so-called damped Ly$\\alpha$ (hereafter DLA) systems are characterized by a very large H~{\\sc i} column density ($N$(H~{\\sc i})~$>$~2$\\times$10$^{20}$ ~cm$^{-2}$), similar to what is usually observed through local spiral disks. The case for these systems to be produced by proto-galactic disks is supported by the fact that the cosmological density of gas associated with these systems is of the same order of magnitude as the cosmological density of stars at present epochs (Wolfe 1996). The presence of heavy elements ($Z \\sim 1/10 ~ Z_\\odot$) and the redshift evolution of metallicity suggest ongoing star formation activities in these systems (Lu et al. 1996, Pettini et al. 1996, 1997). Moreover, strong metal line systems have been demonstrated to be associated with galaxies at low and intermediate $z$ (e.g.  Bergeron \\& Boiss\\'e 1991). It has also been shown that the profiles of the lines arising in the neutral gas show evidence for rotation (Wolfe 1996, Prochaska \\& Wolfe 1997). Whether these arguments are enough to demonstrate that DLA systems arise in large disks is a matter of debate however. Indeed simulations have shown that the progenitors of present day disks of galaxies could look like an aggregate of well separated dense clumps at high redshift. The kinematics could be explained by relative motions of the clumps with very little rotation (Haehnelt et al. 1997, Ledoux et al. 1998). Moreover, using {\\sl HST} high spatial resolution images of the field of seven quasars whose spectra contain DLA lines at intermediate redshifts (0.4~$<$~$z$~$<$~1), Le~Brun et al. (1997) show that, in all cases, at least one galaxy candidate is present within 4~arcsec from the quasar. There is no dominant morphological type in their sample: three candidates are spiral galaxies, three are compact objects and two are amorphous low surface brightness galaxies. Therefore, although the nature of the DLA systems is unclear they trace the densest regions of the Universe where star formation occurs.\\par \\subsection{Molecular hydrogen} \\begin{figure} \\epsfysize=15cm % \\epsfxsize=11cm % \\hspace{1.5cm}\\epsfbox{q0528mol.ps} % \\caption[h]{Fit result for a few rotational transitions of the H$_2$ Lyman absorption bands in the $z_{\\rm abs}$~=~2.8112 system toward PKS~0528--250. The spectrum has been obtained with the echelle spectrograph CASPEC attached on the ESO 3.6~m at La Silla. The resolution is $R$~=~36000 and the integration time 5~hours.} \\end{figure} It is thus surprising that despite intensive searches, the amount of H$_2$ molecules seems quite low in damped Ly$\\alpha$ systems in contrast to what is observed in our own galaxy. Two detections of H$_2$ molecules in high redshift DLA systems have been reported.  Recently Ge \\& Bechtold (1997) have found strong absorptions in the $z_{\\rm abs}$~=~1.9731 DLA system toward Q~0013--004. They derive $N$(H$_2$)~=~6.9$\\times$ 10$^{19}$~cm$^{-2}$, $b$~=~15~km~s$^{-1}$, $T_{\\rm ex}$~$\\sim$~70~K and $n$(H)~$\\sim$~300~cm$^{-3}$ for a total hydrogen column density $N$(H)~=~6.4$\\times$10$^{20}$~cm$^{-2}$. This system has by far the largest H$_2$ abundance $f$~=~2$N$(H$_2$)/[2$N$(H$_2$)~+~$N$(H~{\\sc i})] $\\sim$~0.22$\\pm$0.05 observed in DLA systems. However the exact number should be confirmed using higher resolution data. Other searches have led to much smaller values or upper limits ($f$~$<$~10$^{-6}$, Black et al. 1987, Chaffee et al. 1988, Levshakov et al. 1992). Table~1 summarizes the caracteristics of damped Ly$\\alpha$ systems that have been searched for molecules.  Levshakov \\& Varshalovich (1985) suggested that molecules could be present toward PKS~0528--250 at a redshift ($z_{\\rm abs}$~=~2.8112), slighly larger than the emission redshift of the quasar. This claim has been confirmed by Foltz et al. (1988) using a 1~\\AA~ resolution spectrum. The latter authors derive $N$(H$_2$)~=~10$^{18}$~cm$^{-2}$, $b$~=~5~km~s$^{-1}$, $T_{\\rm ex}$~=~100~K and log~$N$(H~{\\sc i})~=~21.1$\\pm$0.3. By fitting the damped absorption together with the Ly$\\alpha$ emission from the quasar, M\\o ller \\& Warren (1993) find log~$N$(H~{\\sc i})~=~21.35.  New high resolution data has been recently obtained by Srianand \\& Petitjean (1998). They estimate the column density of H$_2$ molecules $N$(H$_2$)~$\\sim$~6$\\times$10$^{16}$~cm$^{-2}$ and the fractional abundance of H$_2$, $f$~=~5.4$\\times$10$^{-5}$ (see Fig.~1). The excitation temperature derived for different transitions suggests that the kinetic temperature of the cloud is $\\sim$200~K and the density $n$~$\\sim$~1000~cm$^{-3}$. The cloud has therefore a dimension along the line of sight smaller than 1~pc. Since it obscurs the broad-line emission region, its transverse dimension should be larger than 10~pc.  Upper limits are obtained on the column densities of C~{\\sc i} ($<$~10$^{12.7}$~cm$^{-2}$) and CO ($<$~10$^{13.2}$~cm$^{-2}$; $N$(CO)/$N$(H~{\\sc i})~$<$~7$\\times$10$^{-9}$). It is suggested that the ratio $N$(H$_2$)/$N$(C~{\\sc i}) is a useful indicator of the physical conditions in the absorber. Photo-ionization models show that radiation fields with spectra similar to typical AGNs or starbursts are unable to reproduce all the constraints and in particular the surprizingly small $N$(C~{\\sc i})/$N$(H$_2$) and $N$(Mg~{\\sc i})/$N$(H$_2$) ratios. In view of the models explored, the most likely ionizing spectrum is a composite of a UV-\"big bump\" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays. This suggests that the gas is not predominantly ionized by the quasar and that star-formation may occur in the clouds, a conclusion reached as well by Warren \\& M\\o ller (1996) and Ge et al. (1997). Dust is needed to explain the production of molecules in the cloud. The amount of dust is broadly consistent with the [Cr/Zn] abundance determination.  \\par\\noindent \\vspace{0.5cm} % \\par\\noindent \\centerline{\\bf Tab. 1 - H$_2$ molecules in DLA systems} \\begin{table}[h] \\hspace{1.5cm} % \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline Name           & 000-263 & 0013-004 & 0100+130  & 0528-250 & 1331+170 & 1337+113 \\\\ \\hline $z_{\\rm em}$   & 4.110 & 2.084  & 2.681  & 2.770  & 2.081 &  2.919\\\\ $z_{\\rm abs}$  & 3.391 & 1.9731 & 2.309  & 2.811  & 1.776 &  2.796\\\\ $N$(HI) (10$^{21}$~cm$^{-2}$ & 2.0 & 0.64  & 2.5  & 2.2  &  1.5  & 0.80\\\\ $N$(H$_2$) (10$^{16}$~cm$^{-2}$& $<$~0.3  & 6900 & $<$~0.5 &  6 & ...  & $<$~5\\\\ $f_{{\\rm H}2}$(10$^{-4}$)  & $<$~0.03  & 2200 & $<$~0.04 & 0.5  & ... & $<$~1.3\\\\ \\hline \\multicolumn{7}{l}{Levshakov et al. (1992); Ge \\& Bechtold (1997); Srianand \\& Petitjean (1998)} \\\\ \\end{tabular} \\label{tsys} \\end{table} \\begin{figure} \\centerline{ \\psfig{figure=boisse.ps,width=9.cm,height=8.cm,angle=270} } \\caption[h]{ [Zn/H] versus log~$N$(H~{\\sc i}) for 37 damped Ly$\\alpha$ systems; small symbols correspond to $z_{\\rm abs}$~$<$~2.15 and large symbols to $z_{\\rm abs}$~$>$~2.15. The line in the upper right corresponds to $N$(Zn~{\\sc ii}) = 1.4$\\times$10$^{13}$~cm$^{-2}$ or to Galactic material inducing $A_{\\rm V}$~$\\sim$~0.27. The figure is taken from Boiss\\'e et al. (1998).} \\label{boisse} \\end{figure} \\subsection{Is there a bias against detection of H$_2$ molecules ?} The small number of H$_2$ detections in damped systems is intriguing. Indeed in the interstellar medium of our Galaxy, all the clouds with log~$N$(H~{\\sc i})~$>$~21 have log~$N$(H$_2$)~$>$~19 (Jenkins \\& Shaya 1979). Formation of H$_2$ is expected on the surface of dust grains if the gas is cool, dense and mostly neutral, and from the formation of negative hydrogen if the gas is warm and dust free (see e.g. Jenkins \\& Peimbert 1997). Destruction is mainly due to UV photons. The effective photodissociation of H$_2$ takes place in the energy range 11.1--13.6 eV, through Lyman-Werner band line absorption.  In the DLA system toward PKS~0528-250, (i) abundances are of the order of 0.1~$Z_{\\odot}$; (ii) the ratio [Cr/Zn] indicates a depletion factor into dust-grains of the order of half of that in the Galactic ISM; (iii) although it has been shown that the cloud is located at a distance larger than 10~kpc from the quasar, it is still close to it and exposed to its UV flux. Nonetheless, molecular hydrogen is detected. This suggests that indeed, molecular hydrogen should be seen in most of the damped systems. The small number of detections may be explained if observations are biased against the presence of molecules. Indeed it can be speculated that molecules should be found predominantly in gas with a non negligible amount of dust. However the corresponding extinction of the background quasar due to the dust in the damped system could be large enough to drop the quasar out of the sample of quasars that are usually observed for such studies. Boiss\\'e et al. (1998) notice that for the damped systems studied up to now, the larger the H~{\\sc i} column density, the smaller the abundances (see Fig.~\\ref{boisse}). This suggests that the high column density DLA systems detected up to now are those with the smallest metallicities and consequently those with the smallest amount of dust. One way to clear up this problem is to observe a complete sample of quasars (if possible constructed without color selection) and to search for high column density damped systems. ",
        "conclusions": ""
    },
    "9804/astro-ph9804061_arXiv.txt": {
        "abstract": "We have computed stellar evolutionary models for  stars in a mass range characteristic of Cepheid variables  ($3<m/\\Msol<12$) for different metallicities representative of the Galaxy and the Magellanic Clouds populations. The stellar evolution calculations are coupled to a linear non adiabatic stability analysis to get self-consistent mass-period-luminosity relations. The period - luminosity relation as a function of metallicity is analysed and compared to the recent EROS observations in the Magellanic Clouds. The models reproduce the observed width of the instability strips for the SMC and LMC. We determine a statistical P-L relationship, taking into account the evolutionary timescales and a mass distribution given by a Salpeter mass function. Excellent agreement is found with the SMC PL relationship determined by Sasselov et al. (1997). The models reproduce the change of slope in the P-L relationship near $P\\sim 2.5$ days discovered recently by the EROS collaboration (Bauer 1997; Bauer et al. 1998) and thus explain this feature in term of stellar evolution. Some discrepancy, however, remains for the LMC Cepheids.  The models are also in good agreement with Beat Cepheids observed  by the MACHO and EROS collaborations. We show that most of the 1H/2H Beat Cepheids have not yet ignited central helium burning; they are just evolving off the Main Sequence toward the red giant branch. ",
        "introduction": "In Fig. 2a-b we compare  our calculations with observed fundamental mode pulsators in the LMC and SMC in a P - $M_V$ diagram. The EROS ($B_E$, $R_E$) magnitudes were transformed into the Johnson-V magnitude according to Beaulieu et al. (1995). We adopt the extinction E(B-V) = 0.10 for the LMC and  E(B-V) = 0.125 for the SMC with R$_V$ = 3.3 and distance moduli $(m-M)_0$ = 18.5 for the LMC and $(m-M)_0$ = 19.13 for the SMC (cf.  Beaulieu et al. 1997a). Comparison is also made with the Laney \\& Stobie (1994, LS94) data, which extend to longer periods than the EROS data. The transformation of the theoretical quantities (L, $\\te$) into M$_V$ are based on the Allard and Hauschildt (1998) most recent atmosphere models. These models originally developed for M-dwarfs do not extend below log g $<$ 3.5. We therefore use the bolometric corrections at constant log g = 3.5 but take metallicity effects into account. We verified however that gravity effects in the V-band are small.  We determine a theoretical statistical P - $\\mv$ relationship under the form $\\mv = A \\, {\\rm log}P + B$ with a weighted least square fit by minimizing the quantity : $$ Q = \\sum_i \\alpha _i \\big( M_{Vi} - A \\, {\\rm log}P_i - B \\big) ^2 $$ where the summation extends over all the stellar masses. The coefficients $\\alpha_i(m_i, t)$ depend on the mass distribution and on the evolutionary times. We adopt a Salpeter mass function (MF) $dN/dm \\propto m^{-2.35}$ and the time dependence derives directly from  the coupled evolution and pulsation calculations. Although the present day MF in the MC may differ from an initial Salpeter MF, we verified that variations of the slope of the MF between -2 and -4 barely affect the slope of the PL relationship. Such a steep MF favors the lowest mass stars because (i) of the number of stars itself and (ii) of the longer time spent in the instability strip as the mass decreases. The slope of the PL relationship is thus hardly affected by stars with m $\\ge 7 \\msol$.  In Fig. 2a, comparison is made between the SMC observations and the Z=0.004 models with masses m= 3 - 12 $\\msol$. The first crossing unstable models are indicated by open circles but only for m $\\le 4 \\msol$. Although included in the calculations, the first crossing instability phase for m $\\ge 5 \\msol$ is statistically insignificant, since the time spent during this phase is more than 300 times smaller than the time spent by a 4 $\\msol$ in the instability strip (first crossing and blue loop). We thus predict the existence of Cepheids with periods $P \\simle \\,1$ day, as the signature of the first instability strip, but the observation of these objects requires larger statistics. The models agree reasonably well with the observed width of the instability strip, although they do not reach the observed blue edge of the EROS data. The faintest objects  ($M_V \\ge -1.5, {\\rm log} P \\le 0.4$) seem to indicate a change of slope in the PL relationship, becoming steeper for log P $\\simle 0.4$, as discovered in the EROS-2 Cepheid sample and analysed carefully by Bauer (1997) and Bauer et al. (1998). We note that such a trend is observed in the theoretical relation near $\\sim$3 $\\msol$. This change of slope is thus real and stems from stellar evolution, illustrating the reduction of the He blue loop as mass decreases (cf. \\S 2). Note also that the period of the minimum mass undergoing a blue loop ($\\sim 3 \\msol$) is consistent with the faintest observed SMC Cepheids. Finally the  average  slope of our P - $\\mv$ relationship is $A$ =  -2.92, in excellent agreement with the slope derived by S97.  Fig. 2b shows the results for the LMC with Z=0.01 and masses from 3.75 to 12 $\\msol$. As in Fig. 2a, the models reproduce correctly the observed width of the instability strip and the overall agreement with observations is good. We note that the minimum theoretical ``unstable'' mass undergoing a blue loop $m_{min} \\sim 3.75 \\msol$ for Z=0.01 does not correspond to the faintest objects observed, which correspond to $\\sim 4.25 \\msol$. We predict that fundamental pulsators with ${\\rm log} P \\simle 0.3$ days should be in the first crossing instability. The slope of the PL relationship with a minimum mass of $4.25 \\msol$ is $A$ = -2.50, shallower than the one observed in the LMC by S97. In order to test the influence of chemical composition, we have recomputed the whole grid of models with (Z=0.008, Y=0.25) and (Z=0.01, Y=0.28), without any substantial modification of the theoretical slope. % For solar metallicity models, we find also a slope $A$ = -2.55 shallower than the one observed in the Milky Way.  A detailed analysis of the possible sources for such discrepancies is under progress. ",
        "conclusions": ""
    },
    "9804/astro-ph9804127_arXiv.txt": {
        "abstract": "A revision of Stod\\'o\\l kiewicz's Monte--Carlo code is used to simulate evolution of star clusters. The new method treats each {\\it superstar} as a single star and follows the evolution and motion of all individual stellar objects. The first calculations for isolated, equal--mass $N$--body systems with three--body energy generation according to Spitzer's formulae show good agreement with direct $N$--body calculations for $N=2000$, $4096$ and $10000$ particles. The density, velocity, mass distributions, energy generation, number of binaries etc. follow the $N$--body results. Only the number of escapers is slightly too high compared to $N$--body results and there is no level off anisotropy for advanced post--collapse evolution of Monte--Carlo models as is seen in $N$-- body simulations for $N \\leq 2000$. For simulations with $N > 10000$ gravothermal oscillations are clearly visible. The calculations of $N=2000$, $4096$, $10000$, $32000$ and $100000$ models take about $2$, $6$ $20$, $130$ and $2500$ hours, respectively. The Monte--Carlo code is at least $10^5$ times faster than the $N$--body one for $N=32768$ with special--purpose hardware (Makino 1996ab). Thus it becomes possible to run several different models to improve statistical quality of the data and run individual models with $N$ as large as $100000$. The Monte--Carlo scheme can be regarded as a method which lies in the middle between direct $N$--body and Fokker--Planck models and combines most advantages of both methods. ",
        "introduction": "Our knowledge about the stellar content, kinematics, and the influence of the environment on observational features of globular clusters and even richer stellar systems are increasing dramatically (Janes 1991, Djorgovski \\& Meylan 1993, Smith \\& Brodie 1993, Hut \\& Makino 1996, Meylan \\& Heggie 1997). First, observations are reaching the point where segregation of mass within globular clusters can be observed directly and quantitatively. Second, observations have revealed that clusters with dense (collapsed) cores are relatively more concentrated to the galactic center than uncollapsed ones. Thus the influences of the environment and mass spectrum are crucial for cluster evolution. Third, observations give clear evidence that post--collapse globular clusters have bluer cores. This suggests strong influence of dynamical interactions between stars on observational properties of globular clusters. Fourth, recent observations show that many different and fascinating types of binaries and binary remnants are present in abundance in globular clusters. Binaries, in addition to being a diagnostic of the evolutionary status of clusters, are directly involved in the physical processes of energy generation, providing the energy source necessary to stop the core collapse and then drive the core expansion. So, to model the evolution of real stellar systems and make meaningful comparison with observation one has to take into account the complex interactions between stellar evolution, stellar dynamics and the environment. Of course all these demands can be fulfilled by direct $N$--body codes (but even the $N$--body method will have trouble with stellar evolution of binary stars). But they are very time--consuming and they need a special--purpose hardware to be run efficiently (Makino 1996ab). Another possibility is to use a code which is very fast and properly reproduces the standard relaxation process and at the same time provides a clear and unambiguous way of introducing all the physical processes which are important during globular cluster evolution. This task might seem unachievable, but actually this kind of code was in use in the past. Monte--Carlo codes, which use a statistical method of solving the Fokker--Planck equation provide all the necessary flexibility.  They were developed by Spitzer (1975, and references therein) and H\\'enon (1975, and references therein) in the early seventies, and substantially improved by Marchant \\& Shapiro (1980, and references therein) and Stod\\'o\\l kiewicz (1986a, and references therein).  Unfortunately, lack of fast computers with sufficient memory at that time and development of the direct Fokker--Planck and gaseous models contribute to the abandonment of this method. But recent developments in computer hardware,  speed and memory now make it possible to run a Monte--Carlo code efficiently, even on general--purpose workstations. The great advantages of this method, beside of its simplicity and speed, are connected with the inclusion of anisotropy and with the fact that added realism does not slow it down. The Monte--Carlo method can practically cope as easily as the $N$--body method with internal freedom of single and binary stars and external environment, with one exception, a stellar system must be spherically symmetric.  The Monte--Carlo code can have another possible use. Despite the simplified nature of continuum models (Fokker--Planck and gaseous models) they will continue for a while to be the most commonly used codes for stellar dynamical evolution. The Monte--Carlo models can be used to optimise physical free parameters and approximations of continuum models to check their validity as it was done in comparison between small $N$--body simulations and continuum ones (Giersz \\& Heggie 1994ab, Giersz \\& Spurzem 1994). This procedure should further increase our confidence in results obtained by Fokker--Planck or gaseous simulations. On the other hand the Monte--Carlo techniques can be incorporated in continuum models to describe the stochastic processes of binary formation, energy generation and movement (Spurzem \\& Giersz 1996, Giersz \\& Spurzem 1997). This, for example, will enable a very detailed investigation of evolution of primordial binaries in evolving background given by an anisotropic gaseous model.  The plan of the paper is as follows. In Section 2 a short review of the `old' and `new' Monte--Carlo methods will be presented. In Section 3 the first results of the `new' Monte--Carlo simulation will be presented. And finally in Section 4 the conclusions and future development of the code will be discussed. ",
        "conclusions": "\\medskip  A successful revision of Stod{\\'o}{\\l}kiewicz's Monte--Carlo code was presented. The updated method treats each {\\it superstar} as a single star and follows the evolution and motion of all individual stellar objects. This improvement was possible thanks to the recent developments in computer hardware and computer speed. Two essential changes was added to the original Monte--Carlo code. Firstly, the procedure which deal with problems of radial velocity determination after the system rearrangement (changes of mechanical energy of the stars due to changes of mass distribution) was slightly changed. This assures better energy conservation. Secondly, the new procedure which deals with star escapers was added. This practically resolves the problem with too high escape rate observed in Monte--Carlo simulations. The Monte--Carlo scheme presented here (as previous Monte--Carlo schemes) takes full advantage of the undisputed physical knowledge on the secular evolution of (spherical) star clusters as inferred from continuum model simulations. Additionally it describes in a proper way the graininess of the gravitational field and the stochasticity of the real $N$--body systems. This does not include any additional physical approximations or assumptions which are common in Fokker--Planck and gas models (e.g. conductivity or isotropic distribution function for field stars). From that respect Monte--Carlo scheme can be regarded as a method which lies in the middle between direct $N$--body and Fokker--Planck models and combines most advantages of the both methods.  The first calculations for equal--mass $N$--body systems with three--body energy generation according to Spitzer's formulae show good agreement with direct $N$--body calculations for $N=2000$, $4096$ and $10000$ particles. The density, velocity, mass distributions, energy generation, number of binaries etc. follow the $N$--body results. Only the number of escapers is slightly too high compared to $N$--body results (but this can be resolved by the time--dependent shift of the escape rate) and there is no level off anisotropy for advanced post--collapse evolution of Monte--Carlo models as is seen in $N$--body simulations for $N \\leq 2000$. For simulations with $N > 10000$ gravothermal oscillations are clearly visible. This is the first unambiguous detection of gravothermal oscillations in Monte--Carlo simulations. Moreover, this is a first unambiguous detection of gravothermal oscillations for stochastic $N$--body system with $N$ as large as $100000$.  The speed of the new code makes it possible to run individual models with $N$ as large as $100000$ and also enables, in an unambiguous way, the inclusion of several different physical processes which operate during different stages of evolution of real globular clusters.  The new Monte--Carlo code described in this paper is seen as a first step towards realistic models of globular clusters. Several important physical processes have to be included to make the simulations of the stellar systems more realistic. The final code will contain the following physical processes: {\\bf (1)} formation of binaries due to dynamical and tidal interactions, {\\bf (2)} primordial binaries, {\\bf (3)} stellar evolution, {\\bf (4)} tidal field of Galaxy and tidal shocks connected with crossing the galactic plane and with large molecular clouds, {\\bf (5)} collisions between stars, {\\bf (6)} interactions between binaries and stars and between binaries themselves, improving the presently used scattering cross-sections for binary hardening.  In the first stage all processes connected with interactions between objects were modelled using analytical cross sections available in the literature. This allowed the code to be tested, and made possible comparison with continuum models.  In the next stage interactions between groups of three and four stars will be modelled by numerical integrations of their orbits (the first attempts are tested now). If during the integration the distance between two or more stars becomes smaller than the sum of their radii then a physical collision takes place. This more realistic approach ensures that processes of energy generation  (the most important factor in the dynamical evolution of globular clusters) will be modelled more closely.  The final stage will be the inclusion of detailed 3--D hydrodynamical modelling of collisions between stars. This will be done by use of Smooth Particle Hydrodynamics (SPH) for a limited number of particles per star (a few hundred).  This will allow close comparison between numerical models and observations of real globular clusters. I refer here to observations of various, peculiar objects like blue stragglers and milliseconds pulsars, which can be formed during collisions and encounters between stars.  \\bigskip \\bigskip  {\\parindent=0pt {\\bf Acknowledgments} I would like to thank Douglas C. Heggie and Rainer Spurzem for stimulating discussions, comments and suggestions to a draft version of this paper. I also thank Douglas C. Heggie, who made the $N$--body results for $N = 4096$ particles available. This work was supported in part by the Polish National Committee for Scientific Research under grant 2--P304--009-06.}"
    },
    "9804/astro-ph9804194_arXiv.txt": {
        "abstract": "We present a Hubble Space Telescope (HST) study of the nuclear region of the E4 radio galaxy NGC 7052, which has a nuclear disk of dust and gas. The Second Wide Field and Planetary Camera (WFPC2) was used to obtain {\\BB}, {\\VV} and~{\\II} broad-band images and an {\\halnii} narrow-band image. The images yield the stellar surface brightness profile, the optical depth of the dust, and the flux distribution of the ionized gas. The Faint Object Spectrograph (FOS) was used to obtain {\\halnii} spectra at six different positions along the major axis, using a $0.26''$ diameter circular aperture. The emission lines yield the rotation curve of the ionized gas and the radial profile of its velocity dispersion. The observed rotation velocity at $r=0.2''$ from the nucleus is $V = 155 \\pm 17 \\kms$. The Gaussian dispersion of the emission lines increases from $\\sigma \\approx 70\\kms$ at $r=1''$, to $\\sigma \\approx 400\\kms$ on the nucleus.  To interpret the gas kinematics we construct axisymmetric models in which the gas and dust reside in a disk in the equatorial plane of the stellar body, and are viewed at an inclination of $70^{\\circ}$. It is assumed that the gas moves on circular orbits, with an intrinsic velocity dispersion due to turbulence (or otherwise non-gravitational motion). The latter is required to fit the observed increase in the line widths towards the nucleus, and must reach a value in excess of $500 \\kms$ in the central $0.1''$. The circular velocity is calculated from the combined gravitational potential of the stars and a possible nuclear black hole. Models without a black hole predict a rotation curve that is shallower than observed ($V_{\\rm pred} = 92\\kms$ at $r=0.2''$), and are ruled out at $>99$\\% confidence. Models with a black hole of mass $\\Mbh = 3.3^{+2.3}_{-1.3} \\times 10^8 \\Msun$ provide an acceptable fit. The best-fitting model with a black hole adequately reproduces the observed emission line shapes on the nucleus, which have a narrower peak and broader wings than a Gaussian.  NGC 7052 can be added to the list of active galaxies for which HST spectra of a nuclear gas disk provide evidence for the presence of a central black hole. The black hole masses inferred for M87, M84, NGC 6251, NGC 4261 and NGC 7052 span a range of a factor $10$, with NGC 7052 falling on the low end. By contrast, the luminosities of these galaxies are identical to within $\\sim\\!25$\\%. Any relation between black hole mass and luminosity, as suggested by independent arguments, must therefore have a scatter of at least a factor $10$. ",
        "introduction": "Astronomers have been searching for direct evidence for the presence of black holes (BHs) in galactic nuclei for more than two decades. Initially, the only constraints on the central mass distributions of galaxies were obtained from ground-based stellar kinematical observations. More recently, the launch of the Hubble Space Telescope (HST) and the subsequent refurbishment in 1993 have provided an important increase in spatial resolution. Combined with new techniques for data analysis and dynamical modeling this has strengthened the stellar kinematical evidence for BHs in several quiescent galaxies (e.g., Kormendy \\etal 1996a,b; van der Marel \\etal 1997a; Cretton \\& van den Bosch 1998; Gebhardt \\etal 1998). New tools for the detection of BHs were also developed. HST observations of the rotation velocities of nuclear disks of ionized gas provided accurate BH mass determinations for several active galaxies (e.g., Harms \\etal 1994; Ferrarese, Ford \\& Jaffe 1996, 1998; Macchetto \\etal 1997; Bower \\etal 1998), while for other galaxies BHs were detected through VLBI observations of nuclear water maser sources (e.g., Miyoshi \\etal 1995). The case for a BH in our own galaxy improved drastically through measurements of stellar proper motions exceeding $1000 \\kms$ in the central $0.1\\pc$ (Genzel \\etal 1997). There are now a total of 10---20 galaxies for which a nuclear dark mass, most likely a BH, has been convincingly detected. The combined results for these galaxies are summarized and reviewed in, e.g., Kormendy \\& Richstone (1995), Ford \\etal (1998), Ho (1998), Richstone (1998), and van der Marel (1998). This sample is now large enough to study the BH mass distribution in galaxies, which is further constrained by ground-based stellar kinematical observations (Magorrian \\etal 1998), HST photometry (van der Marel 1998) and quasar evolution (e.g., Haehnelt \\etal 1998). Our understanding remains sketchy, but is consistent with a picture in which a majority of galaxies has BHs, and in which the BH mass $\\Mbh$ correlates with the luminosity or mass of the host spheroid.  In this paper we present and analyze HST data for the E4 galaxy NGC 7052. This galaxy is a radio source with a core and jet, but no lobes (Morganti \\etal 1987). Ground-based optical images show a nuclear dust disk aligned with the major axis of the galaxy (Nieto \\etal 1990). The physical properties of this disk were discussed by de Juan, Colina \\& Golombek (1996). In a previous paper (van den Bosch \\& van der Marel 1995; hereafter Paper~I) we presented ground-based narrow-band imaging and long-slit spectroscopy obtained with the 4.2m William Herschel Telescope (WHT). These observations showed that there is also a rotating nuclear disk of ionized gas in NGC 7052.  The gas has a steep central rotation curve, rising to nearly $300 \\kms$ at $1''$ from the center. However, the spatial resolution of the spectra was insufficient to convincingly detect a BH, due in part to the relatively large distance of NGC 7052 ($58.7 \\Mpc$; i.e., $1'' = 284.6 \\pc$). The velocity dispersion of the gas was found to increase from $\\sigma \\approx 70\\kms$ at $1''$ from the center to $\\sigma \\approx 200\\kms$ on the nucleus. We showed that this cannot be the sole result of rotational broadening, which would have predicted double-peaked line profile shapes that are not observed.  Instead, the observed central increase in the line width must be at least partly intrinsic. The ground-based kinematics yield an upper limit of $\\sim 10^9 \\Msun$ on the mass of any possible BH.  To improve the constraints on the presence of a central BH we obtained broad- and narrow-band images of NGC 7052 with the Second Wide Field and Planetary Camera (WFPC2) and spectroscopy with the Faint Object Spectrograph (FOS), both in the context of HST project GO-5848. We discuss the imaging and photometric analysis in Section~\\ref{s:WF}, and the spectroscopy and kinematical analysis in Section~\\ref{s:FOS}. In Section~\\ref{s:dyn} we construct dynamical models to interpret the results. With the high spatial resolution of these data we are able to better constrain the nuclear mass distribution, and we find that NGC 7052 has a BH with mass $\\Mbh = 3.3^{+2.3}_{-1.3} \\times 10^8 \\Msun$. We summarize and discuss our findings in Section~\\ref{s:disc}. Some observational details are presented in an appendix.  We adopt $H_0 = 80 \\kms \\Mpc^{-1}$ throughout this paper. This does not directly influence the data-model comparison for any of our models, but does set the length, mass and luminosity scales of the models in physical units. Specifically, distances, lengths and masses scale as $H_0^{-1}$, while mass-to-light ratios scale as $H_0$. ",
        "conclusions": "\\label{s:disc}  We have presented HST observations of the nuclear gas and dust disk in the E4 radio galaxy NGC 7052. WFPC2 broad- and narrow-band images were used to constrain the stellar surface brightness profile, the optical depth of the dust, and the flux distribution of the ionized gas. We have built axisymmetric models in which the gas and dust reside in the equatorial plane, and in which the gas moves on circular orbits with an additional velocity dispersion due to turbulence (or otherwise non-gravitational motion). These models were used to interpret the ionized gas kinematics inferred from our new FOS spectra and from existing ground-based spectra. The models fit the observed central rotation gradient only if there is a central BH with mass $\\Mbh = 3.3^{+2.3}_{-1.3} \\times 10^8 \\Msun$. Models without a black hole are ruled out at $>99$\\% confidence.  The models provide an adequate fit to the available observations with a minimum number of free parameters. The assumptions that we make are similar to those that have been made in HST studies of other galaxies with nuclear gas disks. In several areas our models are in fact more sophisticated than some of the previous work. In particular: we use our multi-colour photometry in order to constrain the central cusp steepness of the stellar mass distribution; we explicitly take into account the contribution of the axisymmetric stellar mass distribution to the circular velocity of the gas, and we do not assume the rotation field to be purely Keplerian; we explicitly model the convolution with the HST/FOS PSF and the binning over the size of the aperture; we model the full line profile shapes, and fit the widths of the emission lines as well as their mean; and we fit Gaussians to the models as we do the data, to properly take into account the fact that Gaussian fits to lines that may be skewed or have broad wings yield biased estimate of the true moments.  Still, our models remain only an approximation to the true structure of NGC 7052. In particular: the thickness of the gas disk may not be negligible; the mean motion of the gas may not be circular; and the observed rotation curve may not perfectly reflect the intrinsic rotation curve, because of partial absorption of the emission line flux by dust. The limited sky coverage of the FOS spectra prevents a direct check on whether the gas motions in NGC 7052 are indeed circular. However, several consistency checks are available that may have signaled errors in our assumptions; none did. The stellar mass-to-light ratio and systemic velocity inferred with our models from the nuclear gas kinematics agree with those inferred from stellar kinematical measurements outside the region influenced by dust absorption. The best-fitting model for the gas kinematics reproduces the shapes of the emission lines on the nucleus, despite the fact that these shapes were not included as constraints in the fit. These agreements do not rule out a conspiracy of some sort, but they do make it less likely that the observed gas kinematics are the result of vastly non-circular motion, or have been strongly modified by dust absorption. Models of adiabatic BH growth for the stellar surface brightness cusp provide another successful check: the BH mass implied by these models is fully consistent with that inferred from the gas kinematics.  Figure~\\ref{f:allBHs} shows a scatter plot of $\\Mbh$ versus {\\BB}-band spheroid luminosity $L_{B,{\\rm sph}}$ for all galaxies with reasonably secure BH mass determinations (adapted from van der Marel 1998, with the addition of NGC 7052; all for $H_0 = 80 \\kms \\Mpc^{-1}$). There is a trend of increasing $\\Mbh$ with increasing $L_{B,{\\rm sph}}$, although it remains difficult to rule out that systematic biases play some role in this relation (van der Marel 1998). Besides NGC 7052, the other galaxies for which the BH detections are based on kinematical studies of nuclear gas disks with the HST are M87 (Harms \\etal 1994; Macchetto \\etal 1997), M84 (Bower \\etal 1998), NGC 6251 and NGC 4261 (Ferrarese, Ford \\& Jaffe, 1998, 1996). The $\\Mbh$ in these galaxies are $3.2 \\times 10^9$, $1.4 \\times 10^9$, $6.6 \\times 10^8$ and $\\Mbh = 4.9 \\times 10^8 \\Msun$, respectively. NGC 7052 falls at the low end of this range. The five galaxies with BH evidence from nuclear gas disks form a very homogeneous set. Each of these galaxies is a radio source and is morphologically classified as an elliptical. The luminosities are identical to within $\\sim\\!25$\\% ($\\log L_B$ in the range $10.6$---$10.8$ for all five galaxies). By contrast, the black hole masses span a range of a factor $10$. The results for these galaxies therefore show that any relation between $\\Mbh$ and $L_{B,{\\rm sph}}$ must have a scatter of at least a factor $10$, even if the comparison is restricted to galaxies of similar type.  \\placefigure{f:allBHs}"
    },
    "9804/astro-ph9804037_arXiv.txt": {
        "abstract": "Two dimensional realizations of self-consistent models for the ``perfect elliptic disks'' were tested for global stability by gravitational N-body integration.  The family of perfect elliptic disk potentials have two isolating integrals; time independent distribution functions $f(E,I_2)$ which self-consistently reproduce the density distribution can be found numerically, using a modified marching scheme to compute the relative contributions of each member in a library of orbits.  The possible solutions are not unique: for a given ellipticity, the models can have a range of angular momenta.  Here results are presented for cases with minimal angular momentum, hence maximal random motion.  As in previous work, N-body realizations were constructed using a modified quiet start technique to place particles on these orbits uniformly in action-angle space, making the initial conditions as smooth as possible.  The most elliptical models initially showed bending instabilities; by the end of the run they had become slightly rounder.  The most nearly axisymmetric models tended to become more elongated, reminiscent of the radial orbit instability in spherical systems.  Between these extremes, there is a range of axial ratios $0.305 \\lesssim b/a \\lesssim 0.570$ for which the minimum streaming models appear to be stable. ",
        "introduction": "Recent studies of elliptical galaxies and of bulges of spiral galaxies indicate that their figures are likely to be at least slightly triaxial (for reviews see \\cite{bin82}; \\cite{dzf91}; \\cite{bs93}).  Most elliptical galaxies appear to be supported at least in part by anisotropies in the velocity distributions rather than by rapid rotation: see, for example, the work on the dwarf elliptical galaxies NGC 147, 185 and 205 by \\cite{bpn91} and \\cite{hdmp92}.  A class of non-rotating potentials, known as the perfect ellipsoids, has been advanced as a possible model for elliptical galaxies (e.g.  \\cite{dez85}).  In these potentials, the mass density is stratified on concentric, similar ellipsoids, and is non-singular in the center.  Many of the properties of these potentials can be derived analytically; the orbits all have three isolating integrals, and hence properties such as the time-averaged density distribution can be computed exactly.  This simplifies the task of finding {\\it self-consistent models\\/}: time-steady phase-space distribution functions $f({\\bf x}, {\\bf v})$ such that the resulting mass density generates the desired gravitational potential. \\cite{stat87} and \\cite{teub87} have demonstrated that distribution functions for the perfect ellipsoids, and the analogous two-dimensional elliptic disks, can be constructed.  Various sub-families of the axisymmetric perfect ellipsoids have been tested for stability (\\cite{dzs89}; \\cite{ms90}; \\cite{mh91}; \\cite{rdz91}).  Flattened perfect ellipsoids could also be viewed as models for galactic bars.  The only analytical bar models are Freeman's (\\cite{free66a}, \\cite{free66b}, \\cite{free66c}) bars, which are based upon a rotating two dimensional harmonic oscillator potential, and the perfect elliptic disk models, which have no figure rotation. \\cite{tdz87} showed that in the limit of the needle ($b \\rightarrow 0$) the two dimensional perfect elliptic disk is neutrally stable. Prompted by the large streaming velocities seen in barred spiral galaxies, the stability of perfect elliptic disks with maximum angular momentum has already been studied (\\cite{slls94}).  The roundest disks were unstable to spiral mode formation, the most elongated elliptical models were unstable to bending modes, while the models with axial ratio $b/a$ in the range $0.250 \\lesssim b/a \\lesssim 0.570$ appeared stable.  This paper extends that previous work to the study of a set of low angular momentum perfect elliptical disks.  The minimal angular momentum cases allow us to study the ability of internal velocity dispersion to support an elliptic figure, and forms a natural complement to the earlier work as the other bound of the whole class of perfect elliptic disks. As before, we tested for global stability by constructing a discrete, self-consistent model, loading it into an $N$-body integrator and allowing it to evolve. ",
        "conclusions": "In this paper and in LS, we have constructed discrete self-consistent representations of the distribution functions of a range of perfect elliptic disks with minimal and maximal angular momentum.  These models were then integrated forward in time using an $N$-body integrator to see if they were stable.  The nearly axisymmetric and the most elongated models were unstable.  The perfect elliptic disks with moderate axial ratios appear to be stable in both the maximum and minimum streaming cases.  In the maximum streaming case, the nearly axisymmetric models developed spiral and bar instabilities as expected, since their limiting case, a cold axisymmetric disk, is known to be violently unstable to spiral instabilities.  In the minimal angular momentum case, nearly-round disks became more elliptical, in a manner very similar to the radial orbit instability of spherical systems.  This is not too surprising, since the velocity distribution is anisotropic, with the radial velocity dispersion being substantially higher than the tangential dispersion, even in the very nearly axisymmetric models.  This comes about because of the substantial presence of box and marginal loop orbits in the models.  In both angular momentum extremes, the most elongated models developed a bending instability.  The similarity in behavior is not very surprising given the decreasing importance of rotational support with increasing ellipticity in these models.  \\cite{tdz87} have shown that the limiting case of the needle ($b\\rightarrow0$) is neutrally stable to bending, while Merritt \\& Hernquist (1991)\\nocite{mh91} have demonstrated a bending instability in a very prolate (E9) system.  It is thus not surprising that the most elliptic models should develop this instability.  For the minimum angular momentum family, the instabilities change the shape of the disk towards a more moderate ellipticity.  In the nearly axisymmetric disks, as the angular momentum is decreased from a maximum, we expect that the increasing velocity dispersion should help to stabilize against spiral instabilities.  It appears likely that there is a stable region for nearly axisymmetric disks with values of Toomre's (\\cite{t64}) stability parameter $Q$ which lies between the points $Q \\sim 2$ and $Q \\sim 3$ where the velocity dispersion has increased to the point of being able to support the disk against the spiral instability (fig \\ref{fig-angmom}).  As the rotational support becomes negligible and the radial velocity dispersion increases, a radial--orbit instability develops; the disks with lower angular momentum become unstable to elliptical distortions when $T_{\\rm radial} / T_{\\rm tangential} \\gtrsim 1.2$ (after the discussion of \\cite{fp84} and \\cite{p87} for stability of spherical systems).  We expect that there is a range of angular momentum between the two extremes for which the nearly round disks are stable.  The moderately elliptical disks with maximum and minimum angular momentum appear to be stable, so we would anticipate that disks of similar ellipticity and intermediate angular momentum will also be stable.  The stability of the two-dimensional models with moderate ellipticity gives us hope that the three dimensional perfect ellipsoids of intermediate triaxiality (which is probably the appropriate range for elliptical galaxies \\cite{mb81}; \\cite{ddc76}; \\cite{dzf91}), will also be stable.  It is known that some very flattened systems, such as the extreme oblate spheroids constructed from thin short-axis tube orbits (\\cite{ms90}), are unstable, but the simple fact that two longer axes are unequal is not likely to be the cause of further trouble.  The three dimensional extension of this work will be interesting to see in light of the work of \\cite{app92} showing that three dimensional systems with a small amount of rotational streaming are unstable to a tumbling bar instability, both when the models have largely radial orbits and when the orbits are mostly circular.  The techniques developed in this work have laid the foundation for investigating the stability of three dimensional perfect ellipsoids, and indeed of any integrable potential.  The methods for choosing orbits, whether simulated annealing (as in LS) or the marching scheme of ZHS, can be easily expanded to take account of a variety of possible cost terms related to angular momentum or line of sight velocities (e.g. \\cite{rix97}).  The procedure of LS for generating a quiet start which minimizes random noise due to particle discreteness, can be carried over to any integrable system.  This is potentially most useful in $N$-body studies which attempt to measure the growth rate of instabilities, because the detection of instabilities which are still in the linear regime is limited by particle noise.  For example, \\cite{s91} found that his linear stability theory was consistent with the results of $N$-body simulations for highly unstable spherical systems, but predicted slow growing instabilities which could not be seen in the simulations because of particle noise.  \\cite{app90} have constructed an analytic potential--smoothing integration technique which decreases the $\\sqrt{N}$ noise associated with binning and softening in $N$-body codes, and permits better examination of the linear growth regime. Their method would also benefit from a quiet start, because the particle discreteness then makes a larger relative contribution to the noise."
    },
    "9804/astro-ph9804171_arXiv.txt": {
        "abstract": "In this letter we present an idea which reconciles a homogeneous and isotropic Friedmann universe with a fractal distribution of galaxies. We use two observational facts: The flat rotation curves of galaxies and the (still debated) fractal distribution of galaxies with fractal dimension $D=2$. Our idea can also be interpreted as a redefinition of the notion of bias.\\\\  {\\bf Key Words:} Large scale structure of universe-cosmology: theory-cosmology:dark matter galaxies:general ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/hep-ph9804444_arXiv.txt": {
        "abstract": "The process of photon splitting $\\gamma \\to \\gamma \\gamma$ in a strong magnetic field is investigated both below and above the pair creation threshold. Contrary to the statement by Baier et al., the ``allowed'' channel $\\alw$ is shown not to be a comprehensive description of splitting in the strong field because the ``forbidden'' channel $\\frb$ is also essential. The partial amplitudes and the splitting probabilities are calculated taking account of the photon dispersion and large radiative corrections near the resonance. \\\\\\\\ PACS numbers: 12.20.Ds, 95.30.Cq, 98.70.Rz ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804165_arXiv.txt": {
        "abstract": "Images and longslit, echelle spectra of the \\Ha emission from 14 dwarf galaxies and M82 have been used to identify expanding shells of ionized gas. Supershells  (radius $>~300$~pc) are found in 12 of the dwarfs. The measured shell sizes and expansion speeds constrain the ages and power requirements of the bubbles. The dynamical age of the larger bubbles is typically about 10~Myr, and ionized shells older than 20~Myr are rare. An energy equivalent to 100 to 10,000 supernova explosions over this period is needed to drive the shock front that sweeps out the cavity. The current star formation rates are high enough to meet these  power requirements. Many of the shells will breakthrough the surrounding layer of HI supersonically, but the projected expansion speeds are typically less than the lower limits on the escape velocity. Some of the shell material may permanently escape from a few galaxies such as \\n1569. Whether bound to the galaxy or not, these outflows probably play an important role in regulating the star formation rate and are expected to significantly influence the chemical evolution of the galaxies. The shells lift gas out of the disk at rates comparable to,  or even greater than, the current galactic star formation rates. They  will only displace a substantial fraction of the interstellar gas if their duty cycle is much longer than the rotational period of the disk. ",
        "introduction": "The interplay between massive stars and the interstellar medium (ISM) plays a fundamental role in the formation and evolution of galaxies. In addition to ionizing radiation and newly synthesized elements, massive stars deliver kinetic energy and momentum  to the surrounding gas through stellar winds and supernova explosions. Shock waves driven by an ensemble of massive stars may trigger additional star formation  and/or sweep the interstellar gas out of the region actively forming stars (Tenorio-Tagle \\& Bodenheimer 1988). The gas flows create a turbulent pressure which helps support the weight of the ISM (e.g. McKee 1990) and  cavities which apparently enhance the distance ionizing radiation propagates (Hunter \\& Gallagher 1997; Martin 1997).  This feedback from star formation may have a particularly strong influence on the evolution of low mass galaxies.  Owing to their low escape velocity, Larson (1974) suggested that the loss of supernova-heated gas would begin earlier and carry away a larger fraction of their initial mass.  This idea was further developed  by Dekel \\& Silk (1988) who used the supernova feedback to regulate the star formation history of the evolving dwarfs. Their starburst-driven wind models were consistent with the observed mass-metallicity and mass-radius scaling relations of dwarfs when a halo similar to those produced in cold dark matter cosmological simulations was included. Mass loss  has subsequently been proposed to explain a number of peculiarities about dwarf galaxies such as their abundance patterns (Marconi,Matteucci, \\& Tosi 1994) and rapid evolution at moderate redshifts (Phillipps \\& Driver 1995; Babul \\& Rees 1992). The ejection of the ISM may not be as easy as previously thought, however. In particular, the rupture of a supershell perpendicular to a galactic disk may vent much of the energy leaving most of the disk intact (DeYoung \\& Heckman 1994).   Observations of dwarf galaxies reveal an environment conducive to the growth of large bubbles.  Their rotation is typically nearly solid body, so  shells are not sheared apart; and metallicities are generally sub-solar so cooling times are longer.  Indeed, small bubbles permeate the star forming regions of the Magellanic Clouds, and a hierarchy of giant shells ($R < 300$~pc) and supergiant shells ($R \\ge 300$~pc) is plainly visible (Davies, Elliot, \\& Meaburn 1976; Meaburn 1980; Kennicutt \\et 1995). The formation of regions like 30~Doradus, which will evolve into a supergiant shell (Chu \\& Kennicutt 1994), may be thought of as the first step in the formation of a galactic outflow. Deep imaging of the ionized gas in other dwarfs yields a plethora of candidate structures for supergiant shells.  Indeed, roughly one out of every four high-surface brightness dwarfs exhibit at least one shell  and/or filaments (Hunter \\et 1993). It is not always obvious, however, which arcs and filaments will show the kinematic signature of an expanding shell (Hunter \\& Gallagher 1990).  The kinematic evidence is mounting that some shells do breakthrough the ambient neutral gas.  In the LMC, for example, the kinematics of many supershells are surprisingly quiescent  compared to the giant shells (Hunter 1994).  Some of these supershells are believed to be the inner ionized surface of cylindrical HI holes (Meaburn 1979; Meaburn 1980; Hunter 1994), which may have formed as a superbubble blew out perpendicular to the galactic plane.  In another Magellanic irregular galaxy, \\n4449, the very large HI hole may be associated with a shell that expanded out of the galactic plane (Hunter \\& Gallagher 1997).  In less luminous galaxies like the blue compact dwarf \\n1705, the expansion of the shell around the central starburst is decidedly non-spherical (Meurer \\et 1992). The kiloparsec scale, expanding shells in amorphous dwarfs (Marlowe \\et 1995) and the faint galaxy IZw18 (Martin 1996) also seem to be elongated in the general direction of the HI minor axis. At issue, however, is whether any of these disk outflows develop into freely flowing winds in which the gas actually escapes from the gravitational potential of the galaxy. Only one member of Marlowe's sample, \\n3955, was a strong wind candidate. The most convincing arguments for actual mass ejection are based on the detection of \\x emitting gas  well above the galactic plane of \\n1569 (Heckman \\et 1995). The association of a hot bubble with the cavity formed by the expanding network of extended \\Ha filaments is reminiscent of the minor axis outflow  from M82 (Bland \\& Tully 1988; Strickland \\et 1996; Shopbell \\et 1997), although it is not yet clear whether the dynamics of these two classes of galactic outflows are completely analogous.   A more extensive kinematic census is desired to assess the frequency of blowout and the amount of mass loss.  This paper presents a catalog of large-scale expanding structures in 14 nearby dwarfs.  Although M82 does not strictly meet the sample selection criteria, it was added to the sample to provide a common galaxy between this study and studies of superwinds from more luminous starbursts (Heckman, Armus, \\& Miley  1990, hereafter HAM). Galaxies were selected from a volume of radius $d \\le 10$~Mpc, right ascension $4 h \\le \\alpha \\le 14~h $, and  declination $\\delta \\ge -35$\\deg.  An effort was made to pick the galaxies with the most intense star formation over a range in absolute luminosity from $M_B = -13.5$ to $M_B = -18.5$. Each radial velocity field was sampled with deep, high-resolution spectra of the \\Ha emission. Additional properties of the galaxies are summarized in Table~\\ref{tab:sam}.  While similar scale shells are found throughout the sample, the net impact on the host galaxy's evolution may be quite varied. Two factors which largely determine the bubble's fate -- i.e. the distribution of the HI and the gravitational  potential -- are not at all uniform across the sample.   Hence, the prospects for mass ejection  are discussed on a galaxy by galaxy basis. The results have interesting applications for the chemical evolution of dwarf galaxies and the regulation of their star formation rate.   This paper is organized as follows. The observations and  data reduction  are described in \\S~\\ref{sec:obs}. Section~\\ref{sec:results} describes the kinematics of the ionized gas, and \\S~\\ref{sec:dynamics} discusses the dynamics of the expanding shells. Rotation curves are sketched in \\S~\\ref{sec:mdot}, and the shell expansion speeds are compared to the escape velocity. Section~\\ref{sec:sum} summarizes the main results. ",
        "conclusions": "\\label{sec:sum}  An extensive set of \\Ha echellograms and images were used to reconstruct the large-scale kinematics of the ionized gas in 14 dwarf galaxies and M82.  Details of the  results for individual galaxies are included in their respective subsections of the  paper, and \\fig~\\ref{fig:weaver} provides a concise summary of the shell expansion speeds and sizes.  The main results regarding the formation of winds in dwarf galaxies are summarized here.  \\begin{itemize}  \\item The formation of supershells must be a common byproduct of massive star formation in dwarf galaxies. Expanding, supergiant ($R > 300$~pc) shells were found in all but two of the galaxies.  This sample was drawn from the population of nearby dwarf galaxies with prominent  arcs and/or extended filaments in their \\Ha emission, and roughly one out of four catalogued dwarf galaxies fits this description (Hunter, Hawley, \\& Gallagher  1993). Indeed, the hierarchical growth of these structures probably began in star forming regions akin to 30~Doradus in the Large Magellanic Cloud (Chu \\& Kennicut 1994). The most powerful  outflows, i.e. \\n1569 and M82, were found to be composed of multiple {\\em cells} whose walls are probably the interface between polar shells. Star formation in the  lowest luminosity galaxies, e.g. IZw18, also generates kiloparsec-scale shells.   \\item Although many of the expanding complexes survive for 10~Myr, none older than $\\sim 20$~Myr were identified.   The lack of shells older than this likely reflects their disruption timescale and provides an indirect measure of the scale height of the ISM. Alternatively, the ionization rate of the shells might drop abruptly on this timescale due to changes in the birthrate of massive stars and/or the illumination geometry.  Although bright, extraplanar HI shells have not been detected in any of the galaxies in this sample, some HI holes in the LMC (Meaburn 1980) and \\n4449 (Hunter \\& Gallagher 1997) are probably relics of expanding supershells. The power requirements of the ionized supershells typically exceed the critical power for supersonic disk breakthrough, so a disruption scenario must be favored for them. The sequence of echellograms stepped across the southern lobe of \\n1569 constrains their deceleration and shows multiple velocity components at least up to 640~pc above the galactic plane. Future observational work must aim to detect the remains of the hot gas and ruptured shell following blowout.  One might speculate that the quiescent filaments in \\n5253 could be fragments of a ruptured shell or clumps of infalling material ejected in a previous wind epoch.  A better census of the local dwarf population would also be helpful for constraining the duty cycle of the winds.    \\item Presuming the shells do rupture, the escape of  hot, X-ray emitting gas from their interiors seems certain. A diffuse, thermal component of the \\x emission has been resolved in three of the galaxies in the sample, but it is  only a significant fraction ($> 10\\%$)  of the interstellar HI mass in M82.  In contrast, much of the interstellar gas swept into the warm ionized shells probably remains bound to the galaxy. The structure of the dark matter halo has been measured in several low surface brightness dwarf galaxies with large HI disks (e.g. Carignan \\& Beaulieu 1989; Meurer \\et 1994) and appears to have a universal structure (Burkert 1995; Navarro, Eke, \\& Frenk 1996). Hence, a conservative approach to mass loss is to assume that the bursting dwarfs are embedded in similar halos.  These dark halo models often do not provide enough mass to explain the HI rotation speed in the inner galaxy, however. Stars and atomic gas can account for essentially all the dynamical mass inside $R(v_{circ})$ in some galaxies like \\n1569, but the fraction varies enormously among the sample members.  Although little CO emission is detected from the dwarf galaxies (e.g. Young \\et 1996), the large uncertainty in the $H_2$ to CO  conversion factor does allow substantial mass contributions from molecular gas (e.g. Maloney \\& Black 1988). A dominant disk in the inner few kiloparsecs has two immediate implications. First, the disk mass contributes significantly to the gravitational acceleration of the kiloparsec-scale shells/disk outflows. Second, the observed turnover in the rotation curve may not be revealing much about the core radius of the halo -- a critical parameter for estimating the escape velocity. The maximum circular velocities in the galaxies do generally appear to be comparable to the expansion velocities of the supershells, so the escape velocities are greater than  the projected shell speed. The expansion speeds along the minor axis of \\n1569 do reach values several times the maximum rotation speed.      \\item The warm shells alone lift gas out of the disk at rates comparable to or greater than the current  galactic star formation rates. The shells transport $10^5$ to $10^6$\\msun\\ of gas over kiloparsec-scale shells in 10~Myr and leave the sound speed high in a large volume of the ISM. This hydrodynamic mixing will be faster than a diffusion process, so the bubbles will clearly alter the chemical evolution of these galaxies. The timescales for blowout are shorter than the evolutionary timescales of most models for Type~Ia supernova progenitors, so the mass loss begins before much of the iron from the burst has been mixed into the ISM. The composition of the ejected material will depend on the duration of the wind and the composition of the ambient ISM, so their impact on the galactic chemical evolution is interwoven with the galactic star formation history.  \\item Although the current kinetic energy in the large expanding structures is only comparable to the binding energy of the ISM in \\n1569, bubble blowout may still extinguish star formation in  particular regions of the other galaxies. If the hot spots percolate across the dwarf irregular galaxies for many rotational periods, then a substantial fraction of the interstellar gas may still be cycled through a halo and/or lost from the galaxy. If the present mass loss rates could be sustained for 6 orbital timescales, for example, then most of the interstellar HI could be removed from the disks of six of the 15 galaxies. This global gas-dynamical feedback will be discussed in the context of the galactic star formation history in a forthcoming paper.   \\end{itemize}"
    },
    "9804/astro-ph9804023_arXiv.txt": {
        "abstract": "Annihilation of high energy, $\\sim 10^{21}$eV, neutrinos on big bang relic neutrinos of $\\sim 1$eV mass, clustered in the Galactic halo or in a nearby galaxy cluster halo, has been suggested to generate, through hadronic $Z$ decay, high energy nucleons and photons which may account for the detected flux of $>10^{20}$eV cosmic-rays. We show that the flux of high energy nucleons and photons produced by this process is dominated by annihilation on the uniform, non-clustered, neutrino background, and that the energy generation rate of $\\sim 10^{21}$eV neutrinos required to account for the detected flux of $>10^{20}$eV particles is $>10^{48}{\\rm erg/ Mpc}^3{\\rm yr}$. This energy generation rate, comparable to the total luminosity of the universe, is $\\sim4$ orders of magnitude larger than the rate of production of high energy nucleons required to account for the flux of $>10^{19}$eV cosmic-rays. Thus, in order for neutrino annihilation to contribute significantly to the detected flux of $>10^{20}$eV cosmic-rays, the existence of a new class of high-energy neutrino sources, likely unrelated to the sources of $>10^{19}$eV cosmic-rays, must be invoked. ",
        "introduction": "The Fly's Eye \\cite{Fly} and AGASA \\cite{AGASA} experiments confirmed the existence of a break in the energy spectrum of high energy cosmic rays at $\\sim5\\times10^{18}{\\rm eV}$, for which evidence existed with weaker statistics in the data of previous experiments (Haverah Park, Yakutsk, Sugar, see e.g. \\cite{Watson} for a review). Fly's Eye data also strengthen the evidence for a change in primary composition from predominantly heavy nuclei below the break to predominantly light nuclei above the break. These features strongly suggest, when coupled with the lack of anisotropy that would be expected for cosmic-rays (CRs) of Galactic origin, that below $\\sim 10^{19}{\\rm eV}$ the CRs are mostly heavy ions of Galactic origin, and that an extra-Galactic component of protons dominates above $\\sim10^{19}{\\rm eV}$. This conclusion is further supported by the fact that the CR energy spectrum is consistent with a cosmological distribution of sources of protons, with injection spectrum $dN/dE\\propto E^{-2.2}$ typically expected for Fermi acceleration \\cite{cosmology}. In particular, there is evidence for the existence of a Greisen-Zatsepin-Kuzmin (GZK) ``cutoff'', i.e. for the suppression of CR flux above $\\sim5\\times10^{19}{\\rm eV}$ expected due to interaction of protons with the microwave background radiation \\cite{GZK}.  The evidence for GZK suppression is strengthened by recent AGASA data \\cite{AGASA1}. In Fig. 1, the CR spectrum reported by the Fly's Eye and the AGASA experiments \\cite{Fly,AGASA1} is compared with the flux expected for a homogeneous cosmological distribution of sources, each generating a power law differential spectrum of high energy protons $dN/dE\\propto E^{-2.2}$ (For the model calculation we have used a flat universe with zero cosmological constant, $H_0=75{\\rm km}\\ {\\rm s}^{-1}$, and time independent energy generation rate per comoving unit volume $5\\times10^{44}{\\rm erg/Mpc}^3{\\rm yr}$; The spectrum is insensitive to the cosmological parameters and to source evolution, since most of the cosmic rays arrive from distances $<1$Gpc \\cite{cosmology}). The deficit in the number of events detected above $5\\times10^{19}{\\rm eV}$, compared to a power-law extrapolation of the flux at lower energy, is consistent with that expected due to a cosmological GZK suppression. However, with current data the ``cutoff'' is detected with only $2\\sigma$ significance \\cite{Tokyo}.  The number of events detected above $10^{20}$eV is consistent with that expected based on the cosmological model presented in Fig. 1 (There is an apparent ``gap'' between the highest and second highest energy events detected by the Fly's Eye \\cite{gap}. However, assuming that the cosmological model is valid, the probability that such an apparent ``gap'' would be observed is $\\sim15\\%$ \\cite{cosmology,gap}). Nevertheless, the detection of $>10^{20}$eV events does pose challenges to most models of CR production. The high energies rule out most of the acceleration mechanisms so far discussed \\cite{Hillas}, and since the distance traveled by such particles must be smaller than $100{\\rm Mpc}$ \\cite{dist} due to their interaction with the micro-wave background, their arrival directions are inconsistent with the position of astrophysical objects, e.g. jets of powerful radio galaxies \\cite{Biermann}, that are likely to produce high energy particles \\cite{obj}.  Cosmological $\\gamma$-ray bursts (GRBs) are likely sources of high-energy CRs, which may account for the CR flux above $10^{19}$eV as well as for the $>10^{20}$eV events \\cite{GRBs}. This model recently gained support form GRB afterglow observations \\cite{AG}. Other models for the production of ultra-high energy CRs were suggested, where the highest energy events are produced by the decay of super-massive elementary particles related to grand unified theories (see, e.g., \\cite{Berezinsky} for recent review). Sources of such particles may be topological defects, left over from a phase transition associated with the symmetry breaking of the grand unified theory \\cite{TD}. While no firm prediction exists of the CR flux in these theories, a generic feature of the super-massive particle decay scenarios is that the injection spectrum is much harder than expected for Fermi acceleration. Therefore, this scenario can account only for the flux of $>10^{20}$eV particles, and can not simultaneously explain the origin of $10^{19}$--$10^{20}$eV CRs.  It has recently been suggested that annihilation of high energy, $\\sim 10^{21}$eV, neutrinos on big bang relic neutrinos of $\\sim 1$eV mass, clustered in the Galactic halo or in a nearby galaxy cluster halo, may generate high energy nucleons and photons which may account for the detected flux of $>10^{20}$eV cosmic-rays \\cite{Weiler}. The existence of $>10^{21}$eV neutrino flux was argued plausible based on the argument that the mechanism producing the observed high energy, $>10^{19}$eV, particles, most likely protons, also produces charged pions of comparable energy, which subsequently decay to produce neutrinos. It was suggested that the generation spectrum extends well beyond $10^{20}$eV, and that while nucleons produced by a distant source, e.g. a powerful radio galaxy, lose their energy interacting with the micro-wave background, high-energy neutrinos propagate without energy losses and may annihilate on relic neutrinos, producing $>10^{20}$eV nucleons and photons at small distances that would allow them to propagate to Earth. In Sec. 2 we derive the energy density of high energy neutrinos required to account for the observed rate of $>10^{20}$eV air showers. The implications of our results are discussed in Sec. 3. ",
        "conclusions": ""
    },
    "9804/hep-ph9804336_arXiv.txt": {
        "abstract": "The broken-symmetry electroweak vacuum is destabilized in the presence of a magnetic field stronger than a critical value. Such magnetic field may be generated in the phase transition and restore the symmetry inside the bubbles. A numerical calculation indicates that the first-order phase transition is delayed but may be completed for a sufficient low value of the Higgs mass unless the magnetic field is extremely high. ",
        "introduction": "It has been found that very strong magnetic fields are capable of destabilizing the electroweak vacuum by forming a vector boson $W^{+}W^{-}$ condensate and restoring the symmetry \\cite{amb}. The required field can only be thought to have existed at the very beginning of the universe and one of the possibilities is that it was generated during the electroweak phase transition \\cite{baym,cheng}. This primordial field may have been subsequently the seed of the present galactic magnetic field \\cite{enq}.  One may wonder whether the restoration of symmetry caused by this strong magnetic field can delay the electroweak phase transition. In particular, if it is of first-order the magnetic field might avoid its completion through the bubble mechanism.  The simplest way to see why a strong magnetic field can destabilize the electroweak vacuum is to consider the energy of a charged spin-one particle interacting with a uniform magnetic field along the 3-axis \\begin{equation} E_{N}^{2}=p_{3}^{2}+m_{0}^{2}+\\left( 2N+1\\right) eB-geB\\quad . \\end{equation}  For the lowest Landau level $N=0$ if the gyromagnetic factor $g$ is $2$ as occurs in the $W$ case, it is clear that the effective mass will become zero for \\begin{equation} B_{c}=\\frac{m_{W}^{2}}{e}\\simeq 10^{24}G\\quad .  \\label{eq2} \\end{equation} This expression is analogous to that of the critical electric field required to create pairs through tunneling.  If one wishes to calculate the decay probability of the vacuum, one must evaluate \\begin{equation} Z=<0|e^{-iHt}|0>=e^{-it\\left( E_{vac}-i\\frac{\\Gamma }{2}\\right) }\\quad . \\end{equation}  In Euclidean metric the one-loop amplitude for a scalar field depends on $ \\det \\left( -D_{E}^{2}+m^{2}\\right) $ with $D_{E\\mu }=\\partial _{\\mu }-ieA_{\\mu }$, being $D_{E4}=iD_{0}$. Using the Schwinger proper time method \\cite{schw} one obtains \\begin{equation} \\ln Z=\\int_{0}^{\\infty }\\frac{ds}{s}tre^{-\\left( -D_{E}^{2}+m^{2}\\right) s}\\quad .  \\label{e2} \\end{equation} For constant electromagnetic fields the trace is known to give the vacuum energy density \\cite{schm} \\begin{equation} \\rho =-\\int_{0}^{\\infty }\\frac{ds}{s}\\frac{e^{-m^{2}s}}{\\left( 4\\pi s\\right) ^{2}}\\left[ \\frac{es\\sqrt{E^{2}-B^{2}}}{\\sin \\left( es\\sqrt{E^{2}-B^{2}} \\right) }-1\\right] \\quad ,  \\label{e1} \\end{equation} where the $-1$ comes from subtracting $\\rho \\left( A=0\\right) $. This integral has a logarithmic divergence for $s=0$ which can be absorbed renormalizing fields and charge \\cite{schw}.  In Eq.(\\ref{e1}) for $E>B$ the integral has poles in the $s$-axis which give origin to an imaginary part corresponding to pair creation. We will be instead interested in the case of $E=0$ and constant magnetic field for which \\begin{equation} \\rho =-\\int_{0}^{\\infty }\\frac{ds}{s}\\frac{e^{-m^{2}s}}{\\left( 4\\pi s\\right) ^{2}}\\left[ \\frac{esB}{\\sinh \\left( esB\\right) }-1\\right] \\end{equation} that has no poles.  For the spin-1 $W_{\\mu }$ case we adopt the view that the only modification to $\\ln Z$ is the interaction of spin with magnetic field $2e\\mathbf{B}\\cdot \\mathbf{s}$ in the exponent of Eq.(\\ref{e2}). Now the trace must be performed on momentum and spin states where the latter involves this added interaction to give \\begin{equation} \\ln Z=\\int_{0}^{\\infty }\\frac{ds}{s}\\left( e^{-2eBs}+e^{2eBs}+1\\right) tre^{-\\left( -D_{E}^{2}+m_{W}^{2}\\right) s}\\quad . \\end{equation} Since the remaining trace is equal to the scalar case, the relevant part of the vacuum energy density is \\begin{equation} \\rho =-\\int_{0}^{\\infty }\\frac{ds}{s}\\frac{e^{-m_{W}^{2}s}}{\\left( 4\\pi s\\right) ^{2}}\\left[ \\frac{esB}{\\sinh \\left( esB\\right) }2\\cosh \\left( 2eBs\\right) \\right] \\quad . \\end{equation} This expression has no poles but diverges for $s\\rightarrow \\infty $ when $ B>B_{c}=m_{W}^{2}/e$ due to the gyromagnetic factor $2$ of the $W$ boson, which would not occur either for $g=1$ or for the $s=1/2$ case. This divergence is an indication of the vacuum instability for large magnetic field. The decay rate should be evaluated in the more realistic situation of $B$ increasing with time, with the consequent generation of an electric field.  In our calculation of the next section we will not take into account the evolution with time of the magnetic field but we will consider that when its value is larger than the critical one in the region of a bubble containing the broken-symmetry vacuum, the bubble will be destroyed. ",
        "conclusions": "We have seen that the highest possible magnetic field together with the most favorable law for having homogeneous field in regions of increasing size might have cosmological consequences through the non-completion of the first-order electroweak transition through a bubble mechanism. Therefore the usual electroweak baryogenesis due to bubble expansion would be affected. But one must notice that also the homogeneous increase of $\\varphi $ can produce a matter-antimatter asymmetry. This is because there will be a baryonic chemical potential related to the time variation of the $CP$ violating phase $\\theta $. The resulting baryonic density will depend on the variation $\\Delta \\theta $ in the interval when the sphalerons are active due to the smallness of $\\varphi $. For a weakly first-order transition, an advantage of this mechanism compared to the bubble one is that the baryonic density would not be erased in the broken phase because here the value of $ \\varphi $ is larger due to the delay of the phase transition. One may remind that this problem is also avoided by the baryogenesis in cosmic strings but paying the price of a suppression factor in the active volume.  However, it is unlikely that such a strong and large size primordial magnetic field has occurred, and for more acceptable fields the effect would be only a small decrease of the temperature for the completion of a first-order transition.  We have studied the influence of the magnetic field on the phase transition using the easiest model, i.e. the standard model and not the MSSM where presumably the first-order phase transition can occur for not too light Higgs mass \\cite{care}. The fact that we obtain the completion of the first-order transition without magnetic field for $m_{H}\\simeq 70GeV$, not far below the experimental bound is probably due to the definition that it occurs when the bubbles touch each other without taking into account their scattering. But we believe that the general conclusions on the magnitude of the effect do not depend on the details of the used electroweak model.  Regarding further developments of this calculation, it would be important to evaluate the vacuum decay rate caused by a time dependent magnetic field in order to consider more carefully its effect on the bubbles instead of taking the simplification of assuming their disappearance as soon as the magnetic field is larger than the critical value.  \\strut  \\strut  \\strut  \\textbf"
    },
    "9804/astro-ph9804090_arXiv.txt": {
        "abstract": "We have computed finite temperature corrections to the electron-hadron scattering cross sections. These are based upon the renormalized electron mass and the modified density of states due to the presence of a background thermal bath. It is found that the electron-hadron thermal transport scattering cross section can be much larger than the zero temperature one. In the case of electron-neutron transport scattering, we find $\\sigma_{ne}(T) / \\sigma_{ne} (T=0) \\simeq 5$ at $T \\simeq 0.1 \\, MeV$.  \\vskip 0.3cm ",
        "introduction": "Finite temperature effects on elementary processes are significant from the point of view of cosmology and astrophysics. The early universe is usually described as a hot gas of particles in nearly thermodynamical equilibrium. Temperature effects enter through the statistical distribution functions. These can renormalize the masses and the wave functions. These renormalized masses and wave functions can then affect scattering processes and decay rates. Several authors \\cite{donoghue} have generalized the electron-mass and wave-function renormalization to all temperatures and densities. Dicus et al. \\cite{dicus} and independently, Cambier et al. \\cite{cps} included the finite temperature effects on weak reaction rates in calculations of standard big-bang nucleosynthesis $(BBN)$. They obtained the corrected light-element abundance and found that the corrections are only of order of a few percent. After that, Saleem \\cite{saleem} included the effects of the electron mass shift at finite temperature on $BBN$ and Baier et al. \\cite{baier} examined the finite temperature radiative corrections to the weak neutron-proton decay rates. More recently, Fornengo et al. \\cite{jwkim} have considered the finite temperature effects on the neutrino decoupling temperature which is important in the evolution of the early universe. In the present work we consider finite temperature corrections to electron-hadron scattering which is important for baryon inhomogeneous cosmologies.  Baryon inhomogeneities might have been produced during the cosmological quark-hadron phase transition in the early universe \\cite{witten}. If such inhomogeneities were present, then the different diffusion lengths for neutrons and protons could lead to the formation of high-baryon density proton-rich regions and low-baryon density neutron-rich regions. The light element nucleosynthesis yields from such regions can differ significantly from those of standard homogeneous big-bang nucleosynthesis \\cite{AHS}. In view of the importance of using light-element yields from the $BBN$ to constrain the baryon-to-photon ratio as well as various cosmological and particle physics theories, such inhomogeneous models must be examined seriously. It is therefore important to quantify the effects of baryon diffusion as accurately as possible.  In this regard Applegate, Hogan and Scherrer (AHS) \\cite{AHS} have calculated the diffusion rate of baryons through the electron-positron plasma in the early universe. Subsequently, several authors used their results in calculations of inhomogeneous $BBN$ \\cite{applegate,mathews}. In $AHS$ it was suggested that the diffusion coefficients could be derived from the mobility of the heavy particles, and that the mobility is determined from the distribution functions of the background plasma and the transport cross section. It is important, therefore, to carefully quantify the values of the distribution functions and the transport cross sections.  However, in all previous baryon diffusion coefficient calculations, vacuum transport scattering cross sections have been used. Therefore, in order to estimate the baryon diffusion coefficients more precisely, in the present work we take into account the finite temperature effects in the calculation of baryon diffusion coefficients at temperatures $\\lsim \\; MeV$.  Specifically, we calculate the transport scattering cross section of elastic electron-hadron scattering at finite temperature. Here we shall treat hadrons as particles which have an internal structure and an anomalous magnetic moment (although we do not have a good field theory for the magnetic moments of protons or neutrons at finite temperature). Also, we assume that their internal structure is not affected by finite temperature since their mass is more than about 1000 times the temperatures of interest.  The plan of the paper is as follows. In Section \\ref{sec:finite}, we discuss how to include finite temperature effects in the calculation. In particular, we will briefly discuss the effective mass of an electron in the MeV temperature range. In Section \\ref{sec:scattering}, we evaluate the electron-hadron transport scattering cross section at finite temperature. Finally, we summarize our results and discuss some astrophysical applications. We shall employ units in which $\\hbar = k_B = c = 1$, except when specific units must be attached to a result. ",
        "conclusions": "We have calculated temperature-dependent electron-hadron transport cross sections. These are important, for example, in the calculation of baryon diffusion coefficients at finite temperature. The major motivation here has been to investigate whether finite temperature effects can significantly change the baryon transport cross section $\\sigma_t$. In this work, we have treated hadrons as particles which have an internal structure and an anomalous magnetic moment. Also, we have assumed that their internal structure is not affected by the finite temperature since their mass is more than about 1000 times the temperature of interest.  Two major features of the finite temperature effects on the light particles have been included in the calculation: (1) finite temperature Dirac spinors which are recast into the form of an effective electron mass ; (2) finite temperature modifications to the phase space distribution of the electrons. We find that, for $m_0 < T$, both $\\sigma_{ne}(T)$ and $\\sigma_{pe}(T)$ approach the ultrarelativistic limit (where the electron mass can be ignored). In the case of electron-proton scattering, we have compared it with the Coulomb scattering cross section at finite temperature. In particular, for the case of electron-neutron transport scattering, we find $\\sigma_{ne}(T) / \\sigma_{ne} (T=0) \\simeq 5$ at $T \\simeq 0.1 \\, MeV$.  In conclusion, the baryon diffusion coefficients which affect baryon inhomogeneities during big-bang nucleosynthesis could be changed significantly by our temperature dependent electron-hadron transport cross sections. Up to the time of weak decoupling ($T \\simeq 1 \\, MeV$) there is little change in the cross sections. However, during the epoch of nucleosynthesis ($T \\lsim 0.2 MeV$) when baryon diffusion is most important, the transport cross sections increase as the temperature decreases. On the other hand, baryon diffusion at low temperature is strongly affected by proton-neutron scattering for which these finite temperature effects are insignificant. Clearly, a study of the effects of these new cross sections on the baryon diffusion coefficients and inhomogeneous primordial nucleosynthesis is desired. These will be the subject of a subsequent paper.  \\vspace{0.5cm} {\\bf Acknowledgements.}  \\noindent The author (ISS) would like to thank P. Marronetti and Prof. S. Rhie for their helpful comments. ISS also acknowledges the support by the Korea Research Foundation (KRF) for the Post-Doctoral Fellowship at University of Notre Dame. This work supported in part by DOE Nuclear Theory Grant DE-FG02-95ER40934."
    },
    "9804/astro-ph9804059_arXiv.txt": {
        "abstract": "To determine the magnification of an extended source caused by gravitational lensing one has to perform a two-dimensional integral over point-source magnifications in general. Since the point-source magnification jumps to an infinite value on caustics, special care is required. For a uniformly bright source, it has been shown earlier that the calculation simplifies if one determines the magnification from the area of the images of the extended source by applying Green's theorem so that one ends up with a one-dimensional integration over the image boundaries. This approach is discussed here in detail, and it is shown that it can be used to yield a robust and efficient method also for limb-darkened sources. It is also shown that the centroid shift can be calculated in a similar way. ",
        "introduction": "For fitting light curves for the ongoing microlensing events, there is a need for robust and efficient methods for calculating the magnification of extended sources, which are not limited to point-lenses. Among the observed events, the presence of binary lenses is a reality (Dominik \\& Hirshfeld~\\cite{MLMC1let},\\cite{MLMC1}; Udalski et al.~\\cite{OGLE7}; Alard et al.~\\cite{DUO2}; Bennett et al.~\\cite{MLMC9}), and planetary events involve a special case of a binary lens. In addition, for some configurations, the light curve for a limb-darkened source will differ significantly from that of a uniformly bright source.  The limb-darkening effect has recently been observed in the galactic microlensing event MACHO 97-BLG-28, which involves both an extended source and a binary lens, by the PLANET collaboration (Albrow et al.~1998a,b\\nocite{M28Planet1}\\nocite{M28Planet2}); the fitting has been done by myself using the algorithm described in this letter.  If one wants to integrate the point-source magnification in two dimensions one has to take special care of the position of the caustics, where the point-source magnification becomes infinite. While this integration can be performed easily for a point-mass lens (e.g. Schneider et al.~\\cite{SEF}, p. 313; Witt \\& Mao~\\cite{WM}; Sahu~\\cite{Sahu}; Dominik~\\cite{DoDiss}), this would be a difficult task for a general lens (e.g. a binary lens), especially at a cusp singularity. In contrast, the area of the images of the extended source and therefore its magnification remains continuous when the source hits a caustic. The determination of the extended source magnification from the boundaries of the image areas has been used to analyze the images of background galaxies behind a cluster of galaxies (Dominik~\\cite{DoDipl}). The image boundaries can be obtained with a contour plot of an implicit function describing the source boundary in the lens plane (Schramm \\& Kayser~\\cite{SK}). This method has been expanded with routines for correcting, testing and finally analyzing the contour line in order to produce an efficient and safe algorithm (Dominik~\\cite{DoAstro}). In that paper, it is noted that it is easy to analyze the images from the contour line data, and an example is given, where quantities such as the area, width, length and curvature of the image have been determined.  Concerning microlensing light curves, it has been noted by Bennett \\& Rhie (\\cite{BenRhie}) that it is advantageous to integrate in the lens plane rather than in the source plane to determine the magnification of an extended source. For uniformly bright sources, Gould \\& Gaucherel (\\cite{GouGau}) proposed applying Green's theorem so that only one integration along the image boundary must be performed rather than two over the image area. This approach is identical to that used earlier (Dominik~\\cite{DoDipl},~\\cite{DoAstro}). The contour plot method is the most convenient way to obtain data points on the image boundary from which the area can be calculated. In Sect.~2, this general approach is described, Sect.~3 gives details for a uniformly bright source, and Sect.~4 shows how this approach can also be used for limb-darkened sources, in which case an easy-to-perform two-dimensional integration remains. In Sect.~5, the calculation of the centroid shift is discussed. ",
        "conclusions": ""
    },
    "9804/astro-ph9804329_arXiv.txt": {
        "abstract": "s{ Knowledge of the constants that describe the current cosmological world model, $H_0$, $t_0$, and the three $\\Omega$s is central to physical cosmology.  Although there is a vast range of suggested tests and existing constraints much of the recent discussion of cosmological constants involves three time variable photospheres: Cepheids, SNIa, and the CMB last scattering surface. These concluding remarks for the Moriond XXXIII meeting are made at a time when many of the established methods have made careful, interesting, statements about the values of various cosmological constants based on data of small random errors. The flood of new data over the next few years will lead to a satisfying increase in the precision of both direct and model dependent estimates of the main cosmological parameters. } ",
        "introduction": "Commenting on the progress being made in estimating the values of the cosmological parameters has perils not unlike those of critiquing a great artwork, perhaps an opera, while it is being staged for the first time. The problem is that there is no working consensus for the values of the cosmological constants. There are rather wide compromise ranges which accommodate most of the derived values. However the values of the cosmological constants at one end of a compromise range are completely incompatible with those at the other, both in physical meaning and in stated measurement error.  An aspect of cosmological parameter determination that is fascinating for any interested scientist is that it continues to reward an integrated view of the entire subject from trigonometric parallaxes to photometric response functions to radio interferometric mapping to X-ray plasma analysis to the outer limits of particle theory, to name but a few. Martin Rees has described Cosmology as the ``Grandest of the Environmental Sciences'' \\cite{rees} which serves to remind us of the difficulty of relating the observations of the universe to the simple cosmological models of interest.  Although the FRW model and an interest in its parameters have been around for a long time, it was not until the 1980's when development of solid state detectors of very low noise and very high quantum efficiency allowed virtually every waveband used by astronomy to greatly increase both the quality and the abundance of data. The objects now examined range from nearby white dwarfs to galaxies at redshifts beyond 5, to precise measurements of the CMB radiation all over the sky.  The following discussion discusses the broad comprise ranges for the basic cosmological parameters. In most cases the data have small random errors.  The problems for all methods come in calibration and model uncertainties.  For instance Cepheids distances have random errors of about 1\\%, nevertheless the total error in the Hubble constant is generally quoted to be at least 10 times larger as a result of potential calibration and systematic errors. Even those apparently large error budgets are hard won from vast efforts to control, measure and remove systematic errors.  The true error ranges of these are continuing to shrink significantly and steadily, such that sometime in the next decade our understanding will undergo a phase transition, hopefully a crystallization not a meltdown. ",
        "conclusions": "The current situation is precisely why the measurement of the cosmological parameters is the primary activity of many astronomers and astrophysicists. The exciting likelihood is that most of the major cosmological constants will be known to a satisfying degree of on the time scale of a decade.  The Moriond meetings provide an ideal format for frank discussion of extremely controversial cosmological issues. I thank the organizers for the splendid job they did in managing to bring us all together and providing a never ending flow of food and stimulus for mind and body."
    },
    "9804/astro-ph9804144_arXiv.txt": {
        "abstract": "The proton burning process $p+p\\rightarrow d +e^+ +\\nu_e$, important for the stellar evolution of main-sequence stars of mass equal to or less than that of the Sun, is computed in effective field theory using chiral perturbation expansion to the next-to-next-to-leading chiral order. This represents a model-independent calculation consistent with low-energy effective theory of QCD comparable in accuracy to the radiative $np$ capture at thermal energy previously calculated by first using very accurate two-nucleon wavefunctions backed up  by an effective field theory technique with a finite cut-off. The result obtained thereby is found to support within  theoretical uncertainties the previous calculation of the same process by Bahcall and his co-workers. ",
        "introduction": "The proton fusion reaction \\be p + p \\rightarrow d + e^+  + \\nu_e \\label{pp}\\ee which plays an important role for stellar evolution and -- as the dominant neutrino source -- for the solar-neutrino problem, has quite a long history of investigation. Indeed the reaction rate of this process (hereafter called the $pp$ rate) was first calculated by Bethe and Critchfield (\\cite{bethe38}). Salpeter (\\cite{salpeter}) recalculated the $pp$ rate using the {\\em effective range approximation} and argued that the relevant nuclear matrix element squared could be estimated with an accuracy of the $\\sim$5 \\% level. (The $pp$ rate itself was subject to much larger uncertainty, $\\sim$20 \\%, because of the limited precision with which the Fermi coupling constant was known at that time.) Bahcall and May (\\cite{bahcall69}) examined the dependence of the $pp$ rate on explicit forms of the two-nucleon wavefunctions generated by two-parameter nuclear potentials of various forms adjusted so as to reproduce the scattering length and effective range (for the $pp$ channel) and the low-energy properties of the deuteron (for the $np$ channel). The $pp$ rate was found to vary by $\\sim 1.5$ \\% corresponding to the changes in the deuteron wavefunction, and by $\\sim 1.2$ \\% due to the change in the $pp$ wavefunction. The most updated work along this line was done by Kamionkowski and Bahcall (\\cite{bahcall94}) employing deuteron wavefunctions obtained from much more accurate potentials such as the Argonne $v_{14}$, $v_{18}$, Urbana $v_{14}$, super-soft-core (SSC) and Reid soft-core potentials. Changes in the $pp$ rates arising from the different potentials were found to be $\\sim 1\\ \\%$. Thus it seems that the presently available calculated $pp$ rate is robust and needs no further scrutiny, the famous solar neutrino problem remaining unresolved from this angle and hence persisting as one of the outstanding unsolved problems in astrophysics (Bahcall \\cite{unsolved}).  There are however two reasons for revisiting this issue. One is that while the calculated $pp$ rate seems to have converged to a ``canonical\"  value given in (Kamionkowski \\& Bahcall \\cite{bahcall94}, hereafter KB), there lingers the unsettling feeling that the strong interaction involved in nuclear physics of the two-nucleon systems is infested with uncontrollable uncertainties associated with model dependence in the treatment, making it difficult to assess the accuracy achieved. Thus it is not unexpected that this canonical value will be -- as has been in the past -- challenged. Indeed it has recently been argued by Ivanov et al. (\\cite{ivanov}) that the nonrelativistic potential models used in the previous works could be seriously in error. They show that in their version of a {\\it relativistic field theory model}, the $pp$ rate comes out to be as big as $2.9$ times the previous estimates.\\footnote{They use  a procedure that seems to disagree with other physical properties of low-energy $pp$ systems, as was pointed out by Bahcall and Kamionkowski (1997). It has also been pointed out by Degl'Innocenti et al. (\\cite{Degl}) that such a large deviation from the value used by KB would be inconsistent with helioseismology in the Sun.} Should their new result turn out to be correct, it would have profound consequences on theories of stellar evolution in general and on the solar neutrino problem in particular. In a nutshell, the issue comes down to whether or not a more general framework such as relativistic field theory would invalidate the calculation made in the traditional nonrelativistic potential models. The claim of the authors in (Ivanov et al. \\cite{ivanov}) is that it indeed does. Our aim is to address this issue using a low-energy effective field theory of QCD that has found a quantitative success in other nuclear processes.  The second reason is really more theoretical, independent of the above important astrophysical issue. Along with the thermal $np$ capture, the proton fusion process is the simplest nuclear process amenable to an accurate calculation -- something rare in hadronic physics -- and it is of interest on its own to test how well a calculation faithful to a ``first-principle approach\" can tackle this problem. In particular, we are interested in checking how accurately the effective field theory approach, found to be stunningly successful for the $np$ capture $n+p\\rightarrow d+\\gamma$, low-energy NN scattering and static properties of the deuteron (Park, Min, \\& Rho \\cite{pmr_PRL}; Park et al. \\cite{pkmr}), fares with the proton fusion problem, the weak interaction sector of the Standard Model.  The strategy we shall adopt here is quite close to that used for the $np$ capture process (Park et al. \\cite{pmr_PRL}). We shall use chiral perturbation theory to the next-to-next-to-leading order (NNLO) in chiral counting; as defined precisely below, this corresponds to $O(Q^3)$ relative to the leading-order term. This is roughly the same order as considered for the $np$ capture. However the relative importance of terms of various chiral orders is somewhat different here.  As we explain later, in the present case, the corrections to the leading order are not suppressed by what is called the ``chiral filter\" \\footnote{The chiral filter phenomenon is explained in detail in (Park, Min, \\& Rho, \\cite{pmr_Report}). Crudely stated, it refers to the general feature that whenever one soft-pion exchange is allowed by kinematics and selection rules, it should give a dominant contribution with higher-order (or shorter-range) terms strongly suppressed. A corollary to this is that whenever one soft-pion exchange is not allowed, all higher-order terms {\\it can be} important, making chiral perturbation calculation generically less powerful.} and so the accuracy with which these can be calculated is not as good as in the $np$ capture case. Even so, using the argument developed in (Park et al. \\cite{pkmr}), we shall suggest that the procedure used here provides a model-independent result in the same sense as in (Park et al. \\cite{pmr_PRL}, Park et al. \\cite{pkmr}).  To streamline the presentation, we first give our result and then discuss (as briefly as possible) how we arrive at it in the rest of the paper. Apart from the meson-exchange contributions which are of order of ${\\cal O}(Q^3)$ and which for the reason mentioned above and further stressed in our concluding section, are the main uncertainty to the order considered, our chiral perturbation theory result in terms of the reduced matrix element $\\Lambda$ defined in (Bahcall \\& May \\cite{bahcall69}) is \\bea \\Lambda_{\\chi PT}^2 = (1 \\pm 0.003) \\times 6.93 \\label{ChPT} \\eea where the uncertainty is due to experimental errors.\\footnote{ Our theoretical uncertainty is, if very conservatively estimated, about $0.1\\ \\%$.} The above result is to be compared with the value obtained by Kamionkowski and Bahcall (KB) \\be \\Lambda_{\\rm KB}^2 = \\left(1 \\pm 0.002^{+0.014}_{-0.009}\\right) \\times 6.92\\, . \\label{bahcallS}\\ee As we shall explain, there are some differences in details between our calculational framework and that of KB, but our final numerical result is in good agreement with that of KB and disagrees with that of Ivanov et al. (\\cite{ivanov}).  The paper is organized as follows. In Section \\ref{2}, our strategy for carrying out a chiral perturbation calculation for two-nucleon systems is outlined. Our approach here is similar to the one used in the previous calculation of the $np$ capture process. We shall sketch a justification of this approach from the standpoint of low-energy effective field theory of QCD (with concrete supporting evidence summarized in Subsection \\ref{cutoff}). Section \\ref{chiralcounting} describes chiral counting of the terms appearing in the relevant weak current. In Section \\ref{wavefunction}, the wavefunctions for the initial $pp$ state and the final $d$ state are specified. Our numerical results are given in Section \\ref{numbers}, and a brief discussion including a comment on the main uncertainty in the calculation is given in Section \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  Following the procedure of chiral perturbation theory proven to be highly successful for the thermal $np$ capture process, we have calculated the $pp$ fusion rate to $\\calO (Q^3)$ in chiral counting relative to the leading single-particle Gamow-Teller matrix element. {\\it To the order considered}, the error involved in the calculation is small, $\\lesssim 1$ \\%. This result, given a justification from a cut-off effective field theory of low-energy QCD as in the case of the $np$ capture (Park et al. \\cite{pkmr}), supports the canonical value of Bahcall et al. and does not support the ``relativistic field theory model\" result of Ivanov et al. (\\cite{ivanov}).  The main caveat in this calculation is in the meson-exchange contribution which comes out to be about 4 \\% when calculated to the chiral order $\\calO (Q^3)$. At the next order, $\\calO (Q^4)$, loops and higher-order counter terms enter, so that there is no reason to believe that they are negligible compared with the $\\calO (Q^3)$ tree contributions. (For instance, it could be lowered to about $1\\sim 2$ \\%  instead of $\\sim 4$ \\% found here). This aspect is different from the case of the $np$ capture where the chiral filter mechanism assures the dominance of the tree-order pion exchange-current contribution. Here the absence of the chiral filter phenomenon {\\it can} allow higher-order (loop) terms to figure equally importantly as the tree-order terms. {}From this viewpoint it is not surprising that a model calculation of the terms of $\\calO (Q^4)$ and higher based on the vector-meson exchange and form factors (Bargholtz \\cite{bar79}) indicates that there can be a considerable suppression of the tree-order correction. Although such a reduction in the exchange-current contribution seems to go in the right direction for beta decays of higher-mass nuclei (Carlson \\cite{carlson}), it is probably unsafe to import the result of such model calculations into our work. To go to $\\calO (Q^4)$ or above in our theory at which heavy mesons and form factors come in, a large number of Feynman graphs of the same chiral order have to be computed on the same footing to assure chiral symmetry, and such a calculation has not been done yet. In the absence of consistent calculations, our attitude is that we should not attach any error estimates on terms not accounted for to the chiral order computed. Calculating the higher-order terms will be left for a future exercise."
    },
    "9804/astro-ph9804002_arXiv.txt": {
        "abstract": "Time-resolved eclipse spectroscopy of the nova-like variable UX~UMa obtained with the HST/FOS on 1994 August and November is analyzed with eclipse mapping techniques to produce spatially resolved spectra of its accretion disc and gas stream as a function of distance from disc centre.  The inner accretion disc is characterized by a blue continuum filled with absorption bands and lines which cross over to emission with increasing disc radius, similar to that reported by Rutten et~al (1994) at optical wavelengths.  The comparison of spatially resolved spectra at different azimuths reveals a significant asymmetry in the disc emission at UV wavelengths, with the disc side closest to the secondary star showing pronounced absorption by an `iron curtain' and a Balmer jump in absorption.  These results suggest the existence of an absorbing ring of cold gas whose density and/or vertical scale increase with disc radius. The spectrum of the infalling gas stream is noticeably different from the disc spectrum at the same radius suggesting that gas overflows through the impact point at disc rim and continues along the stream trajectory, producing distinct emission down to $0.1\\; R_{L1}$.  The spectrum of the uneclipsed light shows prominent emission lines of Ly$\\alpha$, N\\,{\\sc V} $\\lambda 1241$, Si\\,{\\sc IV} $\\lambda 1400$, C\\,{\\sc IV} $\\lambda 1550$, He\\,{\\sc II} $\\lambda 1640$, and Mg\\,{\\sc II} $\\lambda 2800$, and a UV continuum rising towards longer wavelengths. The Balmer jump appears clearly in emission indicating that the uneclipsed light has an important contribution from optically thin gas. The lines and optically thin continuum emission are most probably emitted in a vertically extended disc chromosphere + wind.  The radial temperature profiles of the continuum maps are well described by a steady-state disc model in the inner and intermediate disc regions ($R \\leq 0.3 R_{L1}$). There is evidence of an increase in the mass accretion rate from August to November (from \\.{M}$= 10^{-8.3\\pm 0.1} \\;{\\rm to}\\; 10^{-8.1\\pm 0.1}\\; M_{\\odot} \\; yr^{-1}$), in accordance with the observed increase in brightness.  Since the UX\\,UMa disc seems to be in a high mass accretion, high-viscosity regime in both epochs, this result suggests that the mass transfer rate of UX~UMa varies substantially ($\\simeq 50$ per cent) on time scales of a few months.  It is suggested that the reason for the discrepancies between the prediction of the standard disc model and observations is not an inadequate treatment of radiative transfer in the disc atmosphere, but rather the presence of additional important sources of light in the system besides the accretion disc (e.g., optically thin continuum emission from the disc wind and possible absorption by circumstellar cool gas). ",
        "introduction": "Accretion discs are an important phenomenon in astrophysics, invoked to solve a wide range of astrophysical problems ranging from planetary formation to quasar energetics (Frank, King \\& Raine 1992).  Although considerable effort in both observation and theory has been invested over the past decade, the structure and underlying physics of accretion discs remains poorly understood. Major unsolved problems include the nature of the viscosity mechanism -- responsible for the spiraling inward of the disc material -- (the angular momentum problem), the fate of the kinetic energy expended at the inner edge of the accretion disc (the boundary layer problem), the vertical structure of the disc (the Balmer decrement problem), and the outflow of matter in connection with a disc wind (possibly a solution to, or at least an element of, the boundary layer problem). Progresses in solving these issues has been hampered because most of the previous observational constraints provided only the spectrum of the total light from the disc.  A better understanding of the physics of accretion discs requires spatially-resolved studies.  Cataclysmic Variables (CVs) are mass-exchanging binary systems containing a white dwarf and a late-type star (Warner 1995). If the white dwarf is not strongly magnetized ($B < 10^{6}$~G) an accretion disc is formed.  Accretion discs in non-magnetic CVs cover a range of accretion rates and viscosity states. For example, {\\em dwarf novae} undergo large outbursts ($\\Delta m = 3-5$~mag, typical duration of 5-10 days) which reflects changes in the structure of the discs -- from a cool, optically thin, low viscosity state to a hot, optically thick, high viscosity state -- and which are usually parameterized as a large change in the mass accretion rate ( \\.{M}$= 10^{-11} \\; M_\\odot \\; yr^{-1} \\mapsto 10^{-9} \\; M_\\odot \\; yr^{-1}$.  See, e.g.  Pringle, Verbunt \\& Wade 1986).  On the other hand, {\\em nova-like} variables seem to be permanently in a high viscosity state, presumably as a result of the fact that the accretion rate is always high. Because the nature of the other constituents in these systems -- the white dwarf and the normal star -- are reasonably well understood, and because orbital variations often provide considerable insight into the system geometry, non-magnetic CVs are the ideal laboratories for understanding accretion discs.  Eclipsing systems are particularly useful since the occultation of the accretion disc by the late-type star provides information about the disc's spatial structure through the eclipse shape.  The eclipse mapping method (Horne 1985, 1993; Rutten, van Paradijs \\& Tinbergen 1992; Baptista \\& Steiner 1993) assembles the information contained in the eclipse shape into a map of the disc surface brightness distribution.  When applied to time-resolved spectroscopy through eclipses this technique delivers the spectrum of the disc at any position on its surface.  Information on the radial dependence of the temperature and vertical temperature gradients (for optically thick regions), or temperature, surface density and optical depth (where the disc is optically thin) can be obtained by comparing such spectra with the predictions of models of the vertical disc structure.  The spatial structure of the emission-line regions over the disc can be similarly mapped from data of high spectral resolution. Furthermore, by studying the time-variations in the structure of accretion discs of dwarf novae undergoing outbursts it may be possible to uncover the nature of the (so far unknown) viscosity mechanism which drives accretion discs.  UX~UMa is a well known, bright ($V \\simeq 12.5$) eclipsing nova-like variable with an orbital period of 4.72~hr. Eclipse mapping in broad-bands (Horne 1983; Rutten et~al. 1992) shows that its accretion disc is optically thick and is close to a steady state at a mass accretion rate of $\\simeq 10^{-8} \\; M_\\odot \\; yr^{-1}$.  The broad-band mapping was extended to spectrally-resolved mapping in the optical range by Rutten et al. (1993, 1994).  Their results show that the continuum becomes fainter and redder with disc radius -- reflecting a radial temperature gradient -- and reveal that the Balmer lines are seen in absorption in the inner disc but in emission in the outer disc.  Baptista et al. (1995) performed a similar study using HST data in narrow spectral windows about the C\\,{\\sc IV} 1550 and He\\,{\\sc II} 1640 line regions. This study showed that the UV continuum reasonably follows the $T \\propto R^{-3/4}$ law for steady mass accretion, confirming the results from the optical analysis. The C\\,{\\sc IV} and He\\,{\\sc II} line profiles are dominated by emission from the disc wind. Spatially-resolved spectra reveal that these lines appear as narrow absorption features at disc centre and change with increasing radius to broad emission in the outer disc regions besides showing large uneclipsed components.  This behaviour is similar to that found for the Balmer lines and suggests that these optical lines may also have a wind component.  In this paper, we report on the ultraviolet (UV) and optical mapping of the accretion disc and gas stream of UX~UMa, based on observations made with the {\\it Faint Object Spectrograph} (FOS) on the {\\it Hubble Space Telescope} (HST). The reader is referred to Knigge et~al. (1998a) for an initial description of these observations, with emphasis on the spectral properties of the integrated spectra of the accretion disc, the bright spot and the uneclipsed light.  Sect.\\,\\ref{observa} describes the data and its reduction. The extraction of narrow-band light curves and their analysis with eclipse mapping techniques are described in Sect.\\,\\ref{analise}. Sect.\\,\\ref{results} presents and discusses spatially resolved spectra of the accretion disc and the gas stream region as a function of distance from disc centre, the spectrum of the uneclipsed light, and the radial temperature distribution in the ultraviolet. The possible influence of the assumed eclipse geometry on the results is addressed in Sect.\\,\\ref{geo}.  Sect.\\,\\ref{discuss} discusses the implications of the results in the context of disc atmosphere models. The results are summarized in Sect.\\,\\ref{conclusao}. ",
        "conclusions": "\\label{discuss}  Knigge et~al. (1998a) show that disc models constructed as ensembles of stellar atmospheres provide poor descriptions of the observed integrated spectrum of UX~UMa. The disc model spectra are too blue at ultraviolet wavelengths and overpredict the magnitude of the Balmer jump.  These problems are not new.  The difficulties in fitting integrated spectra of nova-likes and dwarf nova in outburst with disc model spectra have a long history (e.g., Wade 1984, 1988; La Dous 1989; Long et~al. 1991, 1994; Knigge et~al. 1997).  In discussing possible explanations for these problems, Knigge et~al. (1998a) postulated the presence of a significant amount of optically thin material in the system in order to reconcile the disc models with the observed spectrum.  Our spatially resolved study confirms their suggestion by revealing that the integrated spectrum of UX~UMa has indeed a substantial contribution from optically thin emission, most probably associated to the uneclipsed parts of the disc chromosphere + wind. A calculation by Knigge et~al. (1998a) indicated that the addition of an optically thin component with $T= 3 \\times 10^4$\\,K, $n_H= 5 \\times 10^{12}\\; {\\rm cm}^{-3}$, and vertical extension $H= 9.7 \\times 10^9$ cm would be enough to bring the combined disc model plus optically thin emission into good agreement with the observed PRISM spectrum.  The predicted fluxes of their optically thin component raises from $\\simeq 2$ mJy at 2000 \\AA\\ to $\\simeq 5.2$ mJy at 3600 \\AA, being at the level of $\\simeq 3$ mJy at 4500 \\AA\\ -- in good accordance with the fluxes of the uneclipsed component in Fig.\\,\\ref{ffig6}.  For this case, their inferred mass accretion rate is $5 \\times 10^{17}\\; g\\,s^{-1}$ or $10^{-8.1}\\; M_\\odot\\: yr^{-1}$, in excellent agreement with our result.  Thus, the reason for the discrepancies between the prediction of the standard disc model and observations is not an inadequate treatment of radiative transfer in the disc atmosphere (or standard models of vertical structure), but rather the presence of additional important sources of light in the system besides the accretion disc (e.g., optically thin continuum emission from the disc wind and possible absorption by circumstellar cool gas).  Following the same line of reasoning, if disc winds are a common characteristic of all non-magnetic, high state cataclysmic variables, one might expect their disc chromospheres to contribute a non-negligible amount of optically thin emission to the total light of the system. Under this hypothesis, the discrepancy between disc models and the integrated spectrum observed in other non-magnetic, high state systems may be removed by the inclusion of a proper optically thin component to their total light.  These results underscore the importance of spatially resolved studies in disentangling the different components of the integrated spectra of cataclysmic variables. In this particular case, it helped to clarify the situation regarding the apparent discrepancy between disc atmospheres models and the observed spectra."
    },
    "9804/astro-ph9804287_arXiv.txt": {
        "abstract": "Randich and Schmitt [1995, A\\&A 298, 115] found that the coronal activity of solar-type and low mass stars in Praesepe is significantly lower than that of stars in the Hyades cluster.  This result is quite surprising since the Hyades and Praesepe have approximately the same age and metallicity and are often thought to have originated in the same Giant Molecular Cloud complex.  We have carried out several tests in order to find a possible explanation for this result.  We have measured radial velocities  of two groups of Praesepe stars (a dF-dK sample and a dM sample) and have measured H$\\alpha$ as a chromospheric activity index for the dM sample.  Based on analyses of these data, we conclude that the Praesepe catalog used in the X-ray analysis does not contain a significant number of non-members, and thus that membership problems do not seem to be the cause of the  Randich and Schmitt result.  The comparison of the H$\\alpha$ equivalent widths for the M dwarfs in Praesepe with those in the Hyades indicates that, at least for stars in this mass range, the Praesepe stars are as active or more active than their Hyades counterparts.  The similarity of chromospheric emission allows us to reject differences in the rotational velocity distribution as the origin of the dissimilar Lx luminosity functions.  We have also analyzed a few ROSAT PSPC pointings of Praesepe in order to obtain a new and independent estimate of the X-ray luminosities and upper limits for a small sample of Praesepe members. This analysis suggests that the previous ROSAT/PSPC analysis produced slightly optimistic X-ray upper limits; however, the differences between the old and new upper limits are not large enough to explain the dichotomy in the X-ray properties of Praesepe and the Hyades.  Therefore, our examination of the available data does not provide  a clear reason to explain why the X-ray luminosity functions of the two clusters are different. Part of the explanation could be found in the binaries. Speculatively, these clusters could have different orbital period distributions, with more short period binaries among the Hyades, which would show larger coronal activity. ",
        "introduction": "Open clusters play a key role in the understanding of different  time-dependent stellar properties such as the evolution of  rotation, stellar activity,  and   the lithium abundance. The comparison between different clusters which have the same age allows us to prove if this approach is correct or if other effects, such as e.g. a different metal content, are also important. In this report, we examine the Hyades and Praesepe clusters.   During the last 15 years, it has been demonstrated that X-ray emission is a `normal' characteristic of late type stars. As with other stellar properties which depend on rotation (in this case through the dynamo effect, Parker 1955), the emission strength decays with age, as  shown by the comparison between open clusters of different ages  such as the Pleiades (Caillault and Helfand 1985; Micela et al. 1985; Micela et al. 1990; Stauffer et al. 1994), and the Hyades  (Stern et al. 1981; Pye et al. 1994; Stern et al. 1994;  Stern et al. 1995).  The Hyades cluster has been extensively studied at X-ray wavelengths. The ROSAT All-Sky Survey (RASS) detected members down to Log~Lx=1--2$\\times$10$^{28}$ erg~s$^{-1}$, with detection rates of 90\\% for spectral type dG, 40\\% for dK and 30\\% for dM  stars (Stern et al. 1995). They also found that X-ray luminosity functions (XLDF) of K and M-type dwarfs are significantly affected by the presence of a large number of binary systems in the cluster, in the sense that an important fraction of the stars with strong X-ray emission were binaries.  A   study of the X-ray properties of Praesepe has been carried out by Randich and Schmitt (1995). The Hyades and Praesepe have similar age and metallicity although the Hyades are slightly more metal rich. Moreover, their kinetic properties are quite close  (Eggen 1992) and they could have  been  born in the same molecular cloud.  Randich and Schmitt (1995)  presented the results from ROSAT PSPC Raster Scan images in a  4$^\\circ$$\\times$4$^\\circ$ region. Their detection rates for Praesepe were 33\\%, 14\\% and 13\\% for dG, dK and dM stars, respectively. As a consequence, the X-ray luminosity functions of Praesepe in each spectral range are dominated by the upper limits (UL). Since a large fraction of the Praesepe Raster Scan was characterized by a sensitivity similar to that of the Hyades RASS observation, the difference in the detection rates means that the bulk of the Praesepe population is underluminous in X-rays with respect to the Hyades.  The goal of this papers is to try to disentangle this problem, looking for possible reasons of the disparate behavior of these coeval clusters in X-rays.  We present the Praesepe data studied here and the reduction process in Section 2, where in Section 3 we analyze the data and perform a comparison with the  Hyades cluster. Section 4 contains the more important conclusions derived  from this study. ",
        "conclusions": "We have tried to establish the reasons of the different X-ray properties of late type stars in the Hyades and Praesepe. We have studied two different samples of stars: dF-dK Praesepe stars having detected and upper limits for their X-ray luminosities and dM Praesepe stars which have been not detected by ROSAT.  The measured radial velocities for both samples show that contamination by spurious members cannot account for the differences in the level of coronal activity, since all stars (but one) studied here, and presumably most of the stars in the Randich and Schmitt (1995) sample, are real members.  Using simultaneously color-magnitude diagrams and the measured radial velocities, we have discovered new binaries in Praesepe for the dF-dK stars. The comparison of the fraction of binaries in Praesepe and the Hyades shows that it could be slightly different in both clusters. Since the observed  levels of coronal activity, assumed equal sensitivity in the observations, are lower in Praesepe, one would expect a smaller binarity rate in Praesepe than in the Hyades. Moreover, we have shown that the detection rate for the binaries is much higher in the Hyades than in Praesepe. This could be interpreted as an effect of a difference in the distribution of the orbital periods in both clusters.  Finally, the study of the statistical properties of the H$\\alpha$ spectral line for the dM stars in both clusters shows that in fact Praesepe presents higher chromospheric activity for this kind of stars than the Hyades. This result is also surprising, since none of the  Praesepe M dwarfs were detected in X-ray, whereas many of the Hyades M dwarfs are coronally active. For this reason, one possible explanation for the differences in the X-ray properties between both coeval clusters, the existence of different distributions of the rotational velocities, seems unlikely. We have found several Praesepe dM stars which have a remarkable strong H$\\alpha$ emission and very low Lx upper limits, an unexpected situation.  All these data could indicate   the possibility of a difference between the sensitivity of the ROSAT All-Sky Survey for the Hyades and the ROSAT PSPC observations of Praesepe.   However, our re--analysis of several ROSAT pointings shows that the previous assignation of upper limits is essentially correct (although there is a suggestion that the initial estimates for the upper limits were too low by a factor of two). We propose  differences in the orbital period distribution as a partial explanation of the dichotomy of the Lx properties. Extensive studies of different properties which characterize late-type  stars, such as rotational velocities and periods, lithium abundances and additional activity indicators should be made  in a large variety of open clusters in order to have a comprehensive perspective of the evolution of this type of stars. In particular, a similar comparison to that performed here with other clusters of the same age, such as Coma, could also contribute towards an understanding of the differences in the X-ray properties of coeval  clusters. New X--ray data from AXAF or XMM could help to solve this problem."
    },
    "9804/astro-ph9804078_arXiv.txt": {
        "abstract": "The synchrotron reflection scenario recently proposed to explain $\\gamma$-ray flares observed from blazar jets is studied. Our analysis takes into account the angular distribution of the beamed radiation, the finite extent of the scattering region, and light travel-time effects. We compare energy densities and powers for synchrotron, SSC, reflected synchrotron (RSy), and external Compton (EC) scattering processes. If the width of the scattering layer is much larger than $\\Gamma R^\\prime_B$, where $\\Gamma$ and $ R^\\prime_B$ denote the bulk Lorentz factor and comoving-frame radius of the plasma blob, respectively, then the ratio of the RSy and synchrotron energy densities $\\sim 4 \\, \\Gamma^3 n_{\\rm BLR} \\sigma_{\\rm T} R^\\prime_B$, where $n_{\\rm BLR}$ is the mean particle density in the broad line region (BLR). Our results imply that Thomson-thick scattering regions of narrow extent must be present for the synchrotron reflection mechanism to operate effectively. This process seems unlikely to cause flares in lineless BL Lac sources, where X-ray and TeV flares are common and the BLR is thought to be weak or absent. We sketch time profiles of flares for various scenarios, including a model where the blob is energized by sweeping up surrounding material. ",
        "introduction": "More than 50 blazar-type AGNs have been detected with high confidence by EGRET to emit $\\gamma$-rays above 100~MeV (Matox et al. \\markcite{Mattox97}1997). These sources are identified with flat-spectrum radio sources classified as BL-Lac objects or quasars. Many of these objects exhibit variability on all wavelengths, with some of the most rapid variability, on time scales of hours to days, observed at the highest $\\gamma$-ray energies (e.g., Bloom et al. \\markcite{bloom97}1997; Wagner et al. \\markcite{wagner95}1995; Mukherjee et al. \\markcite{mukerjee97}1997).  The large apparent luminosities in combination with the short variability time scales provide evidence for the widely accepted relativistic jet model for AGNs (for recent reviews, see Schlickeiser \\markcite{schlickeiser96}1996 and Hartman et al. \\markcite{Hartman97}1997), according to which the radio--$\\gamma$-ray  emission from blazars is emitted via nonthermal synchrotron radiation and Comptonization of soft photons by energetic particles in relativistic outflows powered by accreting supermassive black holes. Soft photons which are Compton-scattered to produce the $\\gamma$-ray emission include internal synchrotron photons (e.g., Marscher \\& Gear \\markcite{mg85}1985, Maraschi et al. \\markcite{maraschi92}1992, Bloom \\& Marscher \\markcite{bm96}1996) and accretion-disk radiation which enters the jet directly (Dermer \\& Schlickeiser \\markcite{ds93}1993) and after being scattered by surrounding BLR clouds and circumnuclear debris (e.g., Sikora, Begelman \\& Rees \\markcite{sbr94}1994; Blandford \\& Levinson \\markcite{bl95}1995; Dermer, Sturner, \\& Schlickeiser \\markcite{dss97}1997; Protheroe \\& Biermann \\markcite{pb97}1997).  It has recently been proposed (Ghisellini \\& Madau \\markcite{gm96}1996; hereafter GM96) that the beamed synchrotron radiation, after scattering off a cloud near the jet trajectory and reentering  the jet, can be a source of copious soft photons and lead to a pronounced flare of very short duration as the relativistic jet plasma passes through the cloud. Wehrle et al. (\\markcite{Wehrle98}1998) argue that this mechanism might explain the February 1996 $\\gamma$-ray flare observed from 3C 279. A detailed analysis which correctly accounts for causality effects and the finite width of the scattering layer was not performed by GM\\markcite{gm96}96, and such a treatement is required before  blazar flare spectra and light curves can be modeled.  Here we examine this model in more detail for a simple geometry of the BLR in the limiting regime where the nonthermal jet electron distribution does not evolve. In \\S 2  we describe the model. Numerical calculations of the magnetic-field and photon energy densities and synchrotron and Compton powers are presented in \\S 3. Application to blazar flares is made and the effects of blob energization on time profiles of flares are indicated.  We summarize in \\S 4. ",
        "conclusions": "Photon energy densities and radiative powers due to different processes were numerically calculated as a function of the distance of the blob from the accretion disk for a wide range of parameters. Figure 2 shows an example of our series of simulations, using parameters representative of a flat-spectrum radio quasar (FSRQ) and a BL Lac object (BL). Here we let $p = 3$ (FSRQ; solid curves) and $p = 2.7$ (BL; shaded curves). This choice yields fairly flat $\\nu F_{\\nu}$  spectra which can be compared with the peaks of the broadband $\\nu F_{\\nu}$  spectral energy distributions  of blazars. The electron density and blob radius are chosen so that the resulting total luminosities are in accord  with typical values of the apparent luminosities of FSRQs and BLs, and so that $< 100$ GeV $\\gamma$ rays are not absorbed by $\\gamma$-$\\gamma$ pair production on the synchrotron photons intrinsic to the source.  For the generic FSRQ, we assume that the BLR has a radial Thomson depth $\\tau_{\\rm T, BLR} = 0.2$ and occupies a spherical shell located between 0.05 and 0.5~pc from the central engine. For simplicity, we assume that the accretion disk radiates isotropically with a luminosity of $10^{46}$~erg~s$^{-1}$. This choice of parameters gives a Compton power which is $\\sim 10$ times greater than the synchrotron power. Figure 2 shows that although the RSy radiation energy density $u'_{RSy}$ is larger than $u'_{Sy}$ when the blob is located within the BLR, the RSy radiative power is about equal to the SSC power. This is because the backscattered jet synchrotron photons are boosted in energy by a factor of $\\Gamma^2$ relative to the energy of the synchrotron photons in the comoving frame; thus Klein-Nishina effects reduce the radiative power resulting from Compton scattering of RSy radiation more strongly than for the SSC radiation.  In our simulations using plausible parameters compatible with the assumed spherical shell geometry of the BLR, we find that the Compton power due to the synchrotron mirror mechanism is at most comparable to the SSC power.  The efficiency of this process is improved when $\\Gamma \\gg 10$ and a very narrow ($\\Delta r_{\\rm BLR}  \\lesssim r_{\\rm in} / (2\\Gamma^2)$), Thomson-thick BLR is located very far ($r_{\\rm in} \\gg \\Gamma^2 R'_B$) from the central engine. The synchrotron reflection process might therefore operate in BLR clouds which are thought to surround Seyfert AGNs and FSRQs. BLR clouds, as understood through photoionization models (see, e.g., Wandel \\markcite{Wandel97}1997 and references therein), consist of dense ($n\\sim 10^{10}-10^{11}$ cm), Thomson-thick ($\\tau \\sim 1-10$) regions covering a small ($\\sim 10$\\%) fraction of the central engine. For such a model to be feasible, however, the conditions  regarding duty cycle and power outlined by Dermer \\& Chiang (\\markcite{dc98}1998) must be met.  We defer presentation of results in the regime $R^\\prime_B \\gtrsim \\Delta r_{\\rm BLR}/\\Gamma$ to future work; the shell geometry used here is excessively artificial in this limit.  The absence of strong emission lines in X-ray selected BLs and (to a lesser extent) radio-selected BLs suggests that the BLR is considerably more dilute in BLs than in FSRQs, and that BLs have mean Thomson thicknesses $\\tau_{\\rm T,BLR} \\ll 0.1$ (see Scarpa \\& Falomo \\markcite{sf97}1997 and references therein). (On the other hand, the strength of the central ionizing photon source might be much less in BLs than FSRQs.)  Superluminal motion observations also indicate that typical values of $\\Gamma$ for BLs lie in the range between $\\sim 3$ and 7 (see the review by Urry \\& Padovani \\markcite{up95}1995). The ability of the synchrotron reflection process to produce gamma-ray flares therefore seems more difficult in BLs than in FSRQs, yet TeV flares often coincident with X-ray flares have been detected from three BL Lac objects (e.g., Punch et al. \\markcite{punch92}1992, Macomb et al. \\markcite{Macomb95}1995; Catanese et al. \\markcite{Catanese97}1997, \\markcite{Catanese98}1998). For the assumed BL parameters in Figure 2, the synchrotron reflection flare could hardly be detected.  If an accretion disk steadily radiates photons (see B\\\"ottcher \\& Dermer \\markcite{bd95}1995 for a treatment of time-variable disk radiation), then light-travel time effects can be neglected. The accretion disk supplies an abundant supply of soft photons, which can enter the jet directly (ECD) and after being scattered by the BLR (ECC). The comoving photon energy  densities from the ECD process can dominate that from the ECC and synchrotron processes when $z \\lesssim 10^{-2}$~pc, but declines $\\propto z^{-3}$ and $\\propto z^{-2}$ farther out (see Dermer \\& Schlickeiser \\markcite{ds93}1993;  B\\\"ottcher, Mause \\& Schlickeiser \\markcite{bms97}1997). The ECC photon energy density increases slowly with $z$ when $z < r_{\\rm in}$, and begins to decrease for $r_{\\rm in} \\lesssim z \\lesssim r_{\\rm out}$. Outside the BLR, when $z\\gtrsim r_{\\rm out}$, the energy density asymptotically approaches the limiting behavior  $u^\\prime_{\\rm ECC} \\propto z^{-2}$.  When $z \\lesssim r_{\\rm in}$, the ECC process dominates over the SSC process provided that  \\begin{equation} {\\tau_{\\rm T, BLR} \\over r_{\\rm in}^2} \\gtrsim {2 \\over 3} {c \\, B^2 \\tau_B \\over L_D \\, \\Gamma^2} \\> \\big( {p - 1 \\over 3 - p}\\big) \\> {\\gamma_2^{3 - p} - \\gamma_1^{3 - p} \\over \\gamma_1^{1 - p} - \\gamma_2^{1 - p}}, \\end{equation} where $\\tau_B = R^\\prime_B \\, n_{\\rm e,jet} \\, \\sigma_T$ is the Thomson depth of the blob. Using the parameters adopted in Figure 2 but letting $\\tau_{\\rm T, BLR}$, $r_{\\rm in}$, and $L$ vary, we find that the ECC photon energy density dominates the synchrotron photon energy density in the comoving blob frame when $ r_{\\rm in}({\\rm pc}) \\lesssim 0.4 \\> \\tau^{1/2}_{\\rm T, BLR} \\, L_{46}^{1/2}$ and $ r_{\\rm in}({\\rm pc}) \\lesssim 0.1 \\> \\tau^{1/2}_{\\rm T, BLR} \\, L_{44}^{1/2}$ for the FSRQ and BL parameters, respectively, where $L_n = L_{\\rm disk} / (10^n~{\\rm erg \\> s}^{-1})$. It should be noted that Klein-Nishina effects are not included in estimate (6).  The bottom panel of Figure 2 illustrates the time profiles of a flare calculated using the FSRQ parameters for the ECC (thick solid curve)  and RSy (thick dot-dashed curve) processes.  Note that the observer's time element is linearly related to $z$ for  a blob moving with constant velocity. The ECC process gives a fast-rise, power-law-decay--type light curve, and the RSy mechanism gives a  gradual rise of the $\\gamma$-ray flux and a sharp drop as the blob leaves the BLR. A flare produced by the RSy process could be identified by  a rapid decline of $\\gamma$-rays which is not accompanied by a corresponding decrease of the  synchrotron emission. For comparison, we also sketch a flare time  profile produced by the ECD process, and time profiles produced by  sweeping energization of the blob.  In this process, the bulk kinetic  energy of  the outflowing plasmoid is converted into internal nonthermal particle energy by sweeping up BLR material  (see, e.g., Panaitescu \\& M\\'esz\\'aros \\markcite{pm98}1998; Dermer \\& Chiang \\markcite{dc98}1998;  Chiang \\& Dermer \\markcite{Chiang98}1998).  The pair of light solid and dot-dashed curves illustrate ECC and RSy flares, respectively, are modeled assuming that the nonthermal  electron energy is proportional to the amount of swept-up matter which is then added to a nonthermal lepton distribution accelerated at the base of the jet. Blob  deceleration is assumed to be negligible here.  The upper curves of the two pairs are modeled assuming no radiative losses, and the lower curves of the two pairs illustrate the effects of radiative losses on the nonthermal leptons by crudely multiplying the upper curves by a decaying exponential with a $1.5\\times 10^{18}$~cm decay length.  The model RSy light curves, with and without sweeping energization,  are similar to $\\gamma$-ray light curves observed in the 1991 and 1996 flares of 3C 279 (Hartman et al. \\markcite{Hartman96}1996; Wehrle et al. \\markcite{Wehrle98}1998). This may be a consequence, however, of the highly idealized BLR  geometry used in the calculation. More symmetical flaring profiles observed from PKS 0528+134 (Collmar et al. \\markcite{collmar97}1997), PKS 1622-297 (Mattox et al. \\markcite{Mattox97}1997), and PKS 1406-076 (Wagner et al. \\markcite{Wagner95}1995) might be more easily explained by the ECD or ECC processes.  The declines of the X-ray and optical fluxes correlated with the EGRET $\\gamma$-ray fluxes in the February 1996 3C 279 and the 1406-076 flares could, however, rule out the RSy mechanism since the synchrotron component is not directly affected by the reflection process. The inclusion of electron energy evolution and the relaxation of the assumption of a constant velocity blob must be treated to strengthen such conclusions.  Future flare modeling must treat the  passage of the jet through the BLR clouds, which themselves are in  Keplerian motion around the central black hole. The passage of a  jet through such a region will display a complicated signature when monitored by  different telescopes with different sensitivities  and imaging capabilities, which can only be decoded  when full account is taken of the processes considered here.  The efficiency of the different Compton-scattering scenarios,  including the RSy mechanism, in this more realistic system is presently under investigation by the authors."
    },
    "9804/astro-ph9804308_arXiv.txt": {
        "abstract": "We consider the implications of the detection of spiral structure in the accretion disc of the binary IP Pegasi. We use numerical simulations of the development of a disc outburst to construct predicted Doppler tomograms, which are found to be in close agreement with the observations if the spiral pattern arises as a transient feature when the disc expands viscously at the start of the outburst. The good agreement of such viscous disc simulations with the data is consistent with models in which most of the angular momentum transport in the disc originates in internal stresses rather than globally excited waves or shocks. Future detailed observations of the development of transient spiral features offer the potential to measure the dependence of the disc viscosity on the local physical conditions in the disc. ",
        "introduction": "Recent observations of the dwarf nova IP Pegasi provide convincing evidence for spiral structure in the emission from an accretion disc in a binary system (Steeghs, Harlaftis \\& Horne 1997). During an outburst, changes in the profile of spectral lines with binary phase were inverted using the technique of Doppler tomography (Marsh \\& Horne 1988) to reveal a loosely wrapped, two-armed spiral pattern in the disc emission. No such structure is observed in the quiescent disc (Marsh \\& Horne 1990).  These observations provide a potential new constraint on the angular momentum transport processes operating in accretion discs. Two mechanisms are known that can provide a source of viscosity in ionized, non-self-gravitating accretion discs in binary systems; turbulence driven by the non-linear development of the Balbus-Hawley instability (Balbus \\& Hawley 1991; Tout \\& Pringle 1992; Stone et al. 1996; Brandenburg et al. 1996); and spiral waves or shocks driven by the gravitational perturbation of the secondary (Sawada, Matsuda \\& Hachisu 1986; Spruit 1987; Rozyczka \\& Spruit 1989; Savonije, Papaloizou \\& Lin 1994). It is obvious that the second scenario leads to a spiral pattern of disc emission, but even if tidally induced shocks are unimportant for the angular momentum transport budget in the steady-state they might still be observable in outburst, when the enhanced viscosity forces the disc to expand into a region where the strength of the tidal forces is greater (Papaloizou \\& Pringle 1977; Lin \\& Pringle 1976). We note that although it is generally believed that the spiral shock mechanism is inefficient in the relatively cool discs found in cataclysmic variables (Livio 1994; Savonije, Papaloizou \\& Lin 1994), there are considerable theoretical uncertainties in both mechanisms, and additional observational input is highly desirable.  In this Letter, we compare simulations of accretion disc evolution with the observations of IP Peg. Our goal is to test whether viscous disc simulations, which are predicated on the existence of an internal origin for the disc viscosity, are consistent with the strong spiral structure observed in the data. We describe our calculations in Section 2, and present in Section 3 model Doppler tomograms for comparison with the observations. Section 4 summarises our conclusions, and outlines the theoretical expectations for spiral structure in other disc systems. ",
        "conclusions": "In this Letter we have presented a simulation of the evolution of the accretion disc for the parameters of the binary IP Pegasi in outburst. We find as our main result that a spiral pattern is formed in the outer disc during outburst as the enhanced viscous stresses push the disc edge into a region of strong gravitational perturbations from the secondary. The spiral structure obtained in the simulation is two-armed, non-resonant, and much more prominent in outburst as compared to quiescence. Comparing the results with the observations of Steeghs, Harlaftis \\& Horne (1997), we find that there is excellent agreement with the azimuthal extent, velocity range, and asymmetry of the observed pattern.  The observations of IP Pegasi and other cataclysmic variables find no clear evidence for spiral patterns in quiescent discs. This is consistent with theoretical expectations provided that the quiescent discs are cool, viscous, and have not expanded close to the tidal radius. Observations following the decay of spiral structure after the outburst has ended would be valuable in understanding the interplay of these factors. However the current evidence continues to support the conclusion of Savonije, Papaloizou \\& Lin (1994) that angular momentum transport by spiral shocks is insufficiently effective to provide the bulk of the angular momentum transport in the relatively cool, thin discs found in cataclysmic variables. Spiral shocks {\\em are} likely to be important over a wide range of radii in the hotter discs found in X-ray binaries (Owen \\& Blondin 1997), and in the quasi-spherical accretion flows postulated for the advectively dominated regime, although internal sources of viscosity are also likely to be more efficient in those thicker disc geometries.  The current observations can be modelled adequately using a highly simplified three-dimensional model of an outburst caused by a thermal disc instability, in which the only inputs are the change in $\\alpha_{\\rm SS}$ and $c_s$ between quiescence and outburst. Future observations, extending over the rise to outburst and during the decline, may be able to provide stronger constraints on the assumed disc model. In particular, since the spiral pattern arises as a result of the imbalance between internal viscous stresses and well-understood gravitational torques, such observations can probe the variation of the disc viscosity with the local physical conditions in the disc. Eclipsing systems such as IP Peg are particularly promising in this regard as the radial run of quantities such as the effective temperature can readily be derived simultaneously from eclipse mapping.  The detection of spiral structure in the accretion disc of IP Peg, which has relatively feeble outbursts, implies that similar or stronger features may be expected in most dwarf novae. Theoretically, we note that qualitative differences are expected in systems with low mass ratio (roughly, $q < 1/4$). In these binaries both the observations of superhumps in the light curve, and numerical simulations (Murray 1996, 1998), suggest that a strong $m=1$ mode is excited in outburst. These features are {\\em not} stationary in the corotating binary frame, making detailed investigation via Doppler mapping harder. It would also be worthwhile investigating whether spiral structure is observable in magnetic systems such as intermediate polars, where non-axisymmetry might be induced at the {\\em inner} edge of the disc as a result of magnetic torques from the white dwarf.  {\\em Note added:} Godon, Livio \\& Lubow (1998) have recently presented calculations showing that a steady tidally induced spiral pattern does not match the observations of IP Peg. This is consistent with our finding that consideration of the transient behaviour of the disc during outburst is required."
    },
    "9804/astro-ph9804293_arXiv.txt": {
        "abstract": "The Carina Nebula is an extremely bright southern \\HII\\ region embedded in a giant molecular cloud and contains some of the most massive stars known in our Galaxy. We are undertaking a multi-wavelength study of the Carina Nebula in order to examine the detailed kinematics and distribution of the molecular and ionised gas, and to look for further evidence of ongoing star formation. Here we present the results of the initial molecular cloud observations which were made by observing the \\CO\\ emission with the Mopra antenna. The observations reveal the clumpy morphology of the molecular gas, and allow us to identify many interesting regions for follow-up observations. ",
        "introduction": "The Carina \\HII\\ region/molecular cloud complex is an excellent region for studying the interaction of massive stars with their parental Giant Molecular Cloud (GMC). The nebula covers an area of $\\approx$ 4 deg$^{2}$ and is bisected by a prominent V-shaped dark lane. There are over 14 star clusters in this region which have been studied extensively over the past twenty years. For excellent reviews see Feinstein (1995) and Walborn (1995). The most influential clusters of the nebula are the two OB clusters, Tr 14 and Tr 16. These clusters contain numerous O-type stars, including three O3 stars each, making them two of the most massive star clusters in our galaxy. Tr 14 is a compact cluster situated to the north-west of the nebula, adjacent to the western dust lane. Tr 16 is an open cluster centred northwards of the vertex of the dark lane. It contains one of the most massive stars known: $\\eta$ Car. Here we will adopt the popular view (e.g. Tovmassian 1995, Walborn 1995) that Tr 14 and Tr 16 are at a common distance of about 2.2 kpc and that Tr 14 is younger than Tr 16.  Considering the extensive studies on the stellar content of the Carina Nebula, in particular $\\eta$ Car and its surrounding Homunculus nebula, relatively little work has been done on the extended nebula in the last fifteen years. Early radio continuum observations revealed that the nebula contains a large ionised region with two peaks, Car I and Car II (Gardner \\& Morimoto 1968). Higher resolution radio continuum data show that both Car I and Car II are made up of a number of filamentary arcs and rings and are everywhere thermal (Retallack 1983, Whiteoak 1994). Car II is located to the north of $\\eta$ Car and Car I is located towards the western dark lane, just west of Tr 14. The dynamics of the ionised gas in this region have been studied via hydrogen recombination line emission (Gardner et al. 1970, Huchtmeier \\& Day 1975) and H$\\alpha$ and [\\NII] emission observations (Deharveng \\& Maucherat 1975). The results show line splitting towards the Car II region which has been interpreted as an expanding shell of ionised gas.  The dark lanes consist of molecular gas and dust that are associated with the nebula (Dickel 1974). H$_{2}$CO and OH absorption measurements identified two optical depth maxima which were located towards these lanes (Gardner, Dickel \\& Whiteoak 1970, Dickel \\& Wall 1974). Extended far-IR emission is confined there also (Harvey Hoffmann \\& Campbell 1979, Ghosh et al. 1988). There are two main CO emission regions towards the nebula; a northern and southern cloud (de Graauw et al. 1981, Whiteoak \\& Otrupcek 1984). Both regions are part of a much larger GMC which has a projected length of 130 pc and a mass in excess of 5$\\times$10$^{5}$ \\Msun (Grabelsky et al. 1988). The area between the southern and northern CO clouds is centred on the Keyhole Nebula, a dense dark cloud northwest of $\\eta$ Car. Here the molecular gas exists in dense clumps of typical mass 10 \\Msun\\ that are separated both in space and velocity (Cox \\& Bronfman 1995).  The picture used to describe the Carina complex is one in which the massive star clusters, Tr 14 and Tr 16, are interacting strongly with the molecular cloud from which they formed.  It is generally accepted that the photons from Tr 14 and Tr 16 are responsible for the ionised emission of Car I and Car II respectively, and that their strong stellar winds are producing the general expansion of the nebula.  We are undertaking a multi-wavelength study of the Carina Complex in order to study the detailed kinematics and distribution of the molecular and ionised gas and to look for further evidence of ongoing star formation. Here we present the results of initial observations of the \\CO\\ emission. CO emission is thermalised in both low- and high-density gas and therefore is suitable for tracing the overall distribution and velocity structure of the molecular cloud. It also can pinpoint any `CO hot-spots'. These are warm molecular cores where stars could possibly be forming. ",
        "conclusions": "We have presented data from the first stage of an extensive study of the GMC associated with the Carina nebula. The data consists of observations of \\CO\\ emission which have been used to trace the overall GMC as well as pinpoint and CO `hot-spots' or dense regions where stars could possible form. The observations are at a higher resolution than previous studies and reveal the clumpy nature of the northern and southern cloud regions. They also show the positional coincidence between the far-infrared emission and the strong CO emission towards the Car I region. This supports a blister-type model for this region. Further observation of different transitions are being made to better constrain the temperature and density of the molecular gas in this interesting region."
    },
    "9804/gr-qc9804051_arXiv.txt": {
        "abstract": "Linde's proposal of a Euclidean path integral with the ``wrong'' sign of Euclidean action is often identified with the tunneling proposal for the wave function of the universe.  However, the two proposals are in fact quite different.  I illustrate the difference and point out that recent criticism by Hawking and Turok does not apply to the tunneling proposal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804016_arXiv.txt": {
        "abstract": "We report on a near-infrared, long-baseline interferometric search for luminous companions to the star 51~Pegasi conducted with the Palomar Testbed Interferometer.  Our data is completely consistent with a single-star hypothesis.  We find no evidence to suggest a luminous companion to 51~Pegasi, and can exclude a companion brighter than a $\\Delta$K of 4.27 at the 99\\% confidence level for the 4.2-day orbital period indicated by spectroscopic measurements.  This $\\Delta$K corresponds to an upper limit in the companion M$_K$ of 7.30, in turn implying a main-sequence companion mass less than 0.22 M$_{\\sun}$. ",
        "introduction": "The recent inference of a planetary-mass gravitational companion to the star 51~Pegasi (HD 217014) from apparent radial velocity variation by Mayor \\& Queloz (1995) has subjected this otherwise unremarkable star to remarkable scrutiny.  The Mayor and Queloz result was quickly verified by several groups with similar or higher-resolution spectroscopic techniques (c.f.~\\cite{Marcy97}).  However, there has been no other evidence for a companion, e.g.~precision photometric monitoring has failed to show evidence for eclipses (\\cite{Henry97}), and there is a significant lack of x-ray flux from the system compared to binary systems with similar periods (\\cite{Pravdo96}). Further, 51~Peg's G5V spectral classification has become mildly controversial (e.g.~\\cite{Houk95}, who argues for a G2-3V), as has its physical size (e.g.~\\cite{Hatzes97,Henry97}).  A planetary-mass companion in a 4.2 day orbit around a solar-mass 51~Peg would have an orbital semi-major axis of approximately 0.05 AU (\\cite{Marcy97}), slightly more if the companion were more massive. At a distance of 15.4 $\\pm$ 0.2 pc (\\cite{Perryman96}), the approximate maximum primary-companion angular separation would be 3.5 millarcseconds (mas).  Such an angular separation is well below resolution limits for current conventional imaging technology, but is accessible to optical and near-infrared interferometry.  As only the lower mass limit is set by the spectroscopic results, it is possible the companion is significantly more massive -- perhaps even a low-mass star.  We have therefore studied 51~Peg with the Palomar Testbed Interferometer (PTI) in an attempt to detect the putative companion if it is indeed sufficiently luminous.  PTI is a 110m-baseline interferometer operating at K-band (2 -- 2.4 $\\mu$m) located at Palomar Observatory, and described in detail elsewhere (\\cite{Colavita94}).  The minimum PTI fringe spacing is roughly 4 mas at the sky position of 51~Peg, making a (sufficiently) luminous companion readily detectable. ",
        "conclusions": "We find no evidence to suggest that the putative 4.2-day period companion to 51~Peg is detectable in our data; all of the datasets we have analyzed indicate that 51~Peg is at least as stable as our two calibration sources.  The 1997 PTI data on 51~Peg is sufficiently stable that we can place significant limits on $\\Delta$K and consequently M$_K$ of a 4.2-day period companion.  We find upper limits in $\\Delta$K of 4.78, 4.53, and 4.27 for the 4.2-day period companion to 51~Peg at 68\\%, 95\\%, and 99\\% confidence levels respectively.  These $\\Delta$K limits imply companion M$_K$ limits of 7.81, 7.56, and 7.30, corresponding to upper limits on the mass of a putative main sequence companion at 0.17, 0.20, and 0.22 M$_{\\sun}$ at the 68\\%, 95\\%, and 99\\% confidence levels respectively (\\cite{Henry93}).  Our results cannot exclude the possibility of a very low-mass star in a face-on orbit as the 51~Peg companion, but such a star would have to be of spectral type M5V or later."
    },
    "9804/astro-ph9804220_arXiv.txt": {
        "abstract": "We numerically investigate stellar and gas dynamics in star-forming and dissipative galaxy mergers between two disk galaxies  with specific orbital  configurations. We find that violent relaxation combined with gaseous dissipation in galaxy merging transforms two disk galaxies into one S0 galaxy with polar-rings: Both the central S0-like host and the polar-ring component in a polar-ring galaxy are originally  disk galaxies. We also find that morphology of the developed polar-rings reflects both the initial orbit configuration of galaxy merging and the initial mass ratio of the two merger progenitor disk galaxies. Based upon these results, we discuss the origin of the fundamental observational properties of polar-ring galaxies, such as the prevalence of S0 galaxies among polar-ring galaxies, the rarity of polar-ring galaxies among S0 galaxies, the dichotomy between narrow polar-rings and annular ones, shapes of polar-ring warps, and an appreciably larger amount of interstellar gas in the polar-ring component. ",
        "introduction": "Polar-ring galaxies are generally considered to be dynamically peculiar systems in which the outer rings composed of gas and stars are aligned roughly in a perpendicular orientation  with respect to the major axis of the central host galaxies (Schweizer, Whitmore, \\& Rubin 1983; Whitmore et al. 1990; Sackett 1991). A growing number of observational studies have been recently accumulated which can provide valuable information about the origin of these peculiar polar-ring galaxies. Nearly all of the  central host  are morphologically normal S0 galaxies, some of which are confirmed to be rapidly rotating by kinematical studies (Schechter \\& Gunn 1978; Schechter, Ulrich, \\& Boksenberg 1984; Whitmore et al. 1990; Whitmore 1991). Approximately only 0.5 percent of all S0 galaxies have observable polar-rings, which suggests that a particular mechanism is required for the formation of polar-rings in S0  galaxies. The  ring component also shows rapid rotation comparable to that of the main host galaxies, implying that two dynamically different system coexist in these polar-ring galaxies. An appreciable amount of HI gas, which is sometimes comparable to the total mass of the host,  is closely associated with the stellar ring  component (e.g., Shane 1980; Schechter et al. 1984; Richter, Sackett, \\& Sparke 1994; Arnaboldi et al. 1997; Galletta, Sage, \\& Sparke 1997). The morphology of the polar-rings is basically divided into two broad classes (Whitmore 1991): A narrow ring  which is not extended in size  (e.g., ESO 415-G 26) and an annulus which is a disk-like component with  the central part cut out (e.g., NGC 4650a). Peculiar morphology is  observed in some polar-ring galaxies (e.g., the Helix galaxy, NGC 2685,  and double ringed system, ESO 474-G 26), which further implies  considerably complicated physical processes in  polar-ring formation and simultaneously provides a clue to the understanding of the origin of polar-ring galaxies (Sackett 1991). Roughly two-thirds of these polar-rings show  obvious galactic  warps whose shapes look like `integral sign' and/or   `banana'(Whitmore 1991). Statistical studies on the distribution of the angle between the ring component and the central host reveal that these two components strongly prefer to be orthogonal with each  other. These peculiarities both in the kinematics and morphology observed in polar-ring galaxies have attracted a number of  theoretical interests, which are divided basically into two categories: One is the origin of the polar-rings and the other is the nature of dark matter halo surrounding  polar-ring  galaxies. Although there are a large number of important studies addressing the three dimensional shapes of dark matter halo in the galaxies (e.g., Whitmore, Mcelroy, \\& Schweizer 1987; Reshetnikov \\& Combes 1994; Sackett et al. 1994; Combes \\& Arnaboldi 1996), we here restrict ourselves to the mechanisms which would naturally explain the formation of the polar-ring galaxies with spherical haloes.   It is generally believed that the formation of the polar-rings is the results of a `secondary event' involving a pre-existing S0 galaxy (e.g., Steiman-Cameron \\& Durisen 1982; Sparke 1986; Quinn 1991; Rix \\& Katz 1991; Reshetnikov \\& Sotnikova  1997). Specifically, the host S0 galaxy is supposed  to have acquire the material constituting the ring component by capturing the gas during tidal interaction with  neighbor galaxies. The subsequent gravitational interaction combined with the gaseous dissipation then spreads the captured gas and  forms  the polar-rings around the host galaxy. One of the promising models along  this orthodox scenario is the `preferred plane model' in which the differential precession of the rings and the gaseous dissipation cooperate to play a vital role in leading the acquired gas to settle into the stable polar orbit and finally to form the polar-rings (Tohline \\& Osterbrock 1982; Durisen et al. 1983; Schweizer et al 1983). A number of numerical simulations have already confirmed in what physical conditions the polar-rings are more likely to form and continue to exist for a relatively longer time-scale (Habe \\& Ikeuchi 1988; Christodoulou et al. 1992; Katz \\& Rix 1992). Indeed these previously proposed models have provided a potential success in reproducing the polar-rings in S0 galaxies, however, these seem to be incapable of giving sufficiently conclusive and persuasive answers to the following seven  questions on the origin of the polar-ring galaxies (Sackett 1991; Whitmore 1991; Arnaboldi et al. 1997; Galletta et al. 1997): (1) Why are nearly all the central host galaxies morphologically classified as S0 ? (2) Why are  polar-ring galaxies so rare among S0 galaxies ? (3) Why do some polar-ring galaxies have a narrow  ring and some have annuli ? (4) Why are the mass and angular momentum of the ring component comparable or sometimes larger than  those of the host ? (5) Why are the rings so `polar' ? (6) Why some polar-rings have considerably peculiar morphology such as helical and double-ringed shapes ? (7) Why is there an appreciably greater amount of interstellar gas in the polar-ring component ? In particular, (1),  (4), and (7) could not be explained  simply by the previous theoretical  models, implying  either that more elaborated and sophisticated models along the above orthodox  scenario should be considered or that the alternative model should be proposed for the explanation of the above questions.   The purpose of this paper is to explore the origan of polar-ring galaxies and to propound a new mechanism which more naturally  and reasonably explains the aforementioned observational trends of polar-ring galaxies. In the present study, we consider that the dissipative galaxy merging between two disks is a promising mechanism that quite reasonably answers  the above seven questions. Therefore we investigate how the dissipative galaxy merging transforms two disks into one early-type S0 galaxy with polar-ring. Furthermore, we investigate how the orbit configuration of galaxy merging and initial mass ratio of the two progenitor disk galaxies can affect the morphology of polar-rings developed after galaxy merging. In this paper, the galaxy merging with specific orbit configurations and sufficient  amount of gaseous dissipation is demonstrated to play a vital role in forming both the cental S0-like host and the surrounding polar-ring component in polar-ring galaxies. This paper is an extended version of Bekki (1997) in which the basic mechanism of polar-ring S0 galaxy formation in galaxy mergers is briefly summarized. The layout of this paper is as follows.  \\S 2 describes numerical models for dissipative galaxy merging. \\S 3 gives the results obtained in the present study. In \\S 4, we mainly discuss whether or not  the model proposed in the present paper can become a new promising model which naturally and reasonably explains the observational properties of polar-ring S0 galaxies. ",
        "conclusions": "The present numerical study provides a new mechanism by which both the central S0-like host and the ring component in a polar-ring galaxy are simultaneously formed. Although uncertainties of the numerical treatment of gas dynamics and star formation still remain, it appears that our model has succeeded in reproducing $some$ polar-ring galaxies and explaining naturally a number of important observational properties of them. In the proposed model, the formation of polar-ring galaxies is essentially ascribed to the details of dynamics of dissipative galaxy merging with specific  orbital  configurations. Specifically, the central host of a polar-ring galaxy is the galaxy which has been inevitably transformed from a late-type spiral into an early-type S0 galaxy  during the  merging. The ring component, on the other hand, is the `galaxy' which has been dramatically transformed from a late-type spiral into a narrow  ring or annuli  owing to the violent gravitational interaction and gaseous dissipation during  the merging. Although both specific  orbital  configurations and gaseous dissipation in galaxy merging are required for the formation of polar-ring galaxies in the present  model, these constraints also  give natural explanations to observed trends  such as the prevalence of S0 among polar-ring galaxies (e.g., Whitmore 1990), the rarity of the polar-ring galaxies  among S0 galaxies (e.g., Whitmore et al. 1991), and an appreciably larger amount of interstellar gas in polar-rings (e.g., Sackett 1991). Moreover it  is found that the morphology of  polar-rings such as a narrow ring (e.g., ESO 415 - G 26), annular rings (e.g., NGC 4650a), helical rings (e.g., NGC 2658), and double rings (e.g., ESO 474 - G 26) can reflect both the orbital  parameters of galaxy merging and the initial mass ratio of two merger precursor galaxies. Thus, the present study demonstrates that a merger remnant of a gas-rich galaxy merger is one of promising candidates of polar-ring S0 galaxies."
    },
    "9804/astro-ph9804150_arXiv.txt": {
        "abstract": "We present a new, fully covariant and manifestly gauge-invariant expression for the temperature anisotropy in the cosmic microwave background radiation resulting from scalar perturbations. We pay particular attention to gauge issues such as the definition of the temperature perturbation and the placing of the last scattering surface. In the instantaneous recombination approximation, the expression may be integrated up to a Rees-Sciama term for arbitrary matter descriptions in flat, open and closed universes. We discuss the interpretation of our result in the baryon-dominated limit using numerical solutions for conditions on the last scattering surface, and confirm that for adiabatic perturbations the dominant contribution to the anisotropy on intermediate scales (the location of the Doppler peaks) may be understood in terms of the spatial inhomogeneity of the radiation temperature in the baryon rest frame. Finally, we show how this term enters the usual Sachs-Wolfe type calculations (it is rarely seen in such analyses) when subtle gauge effects at the last scattering surface are treated correctly. ",
        "introduction": "The calculation of the primary temperature anisotropy in the cosmic microwave background radiation (CMB) resulting from density perturbations has a long history, beginning with the seminal paper by Sachs and Wolfe~\\cite{sachs67}. Since the original Sachs-Wolfe estimate, a wealth of detailed predictions for the anisotropies expected in various cosmological models have been worked out. The calculations are straightforward in principle, but, like many topics in cosmological perturbation theory, are plagued by subtle gauge issues~\\cite{stoeger95a}.  The problems of gauge-mode solutions to the linear perturbation equations and the gauge-ambiguity of initial conditions can be eliminated by working exclusively with gauge-invariant variables, as in the widely used Bardeen approach~\\cite{bardeen80} and the less well known covariant approach advocated by Ellis and coworkers~\\cite{ellis89a,ellis89b}. However, gauge issues still arise in connection with the definition of the temperature perturbation and the placement of the last scattering surface~\\cite{stoeger95a,ellis-er97}. The latter gauge issues do not arise at first-order in numerical calculations which integrate the Boltzmann equation in a perturbed universe, since the visibility function (which determines the position of the last scattering surface) multiplies first-order variables giving only a second-order error from the use of a zero-order approximation to the visibility~\\cite{LC-scalcmb}. However, this is not always the case in Sachs-Wolfe style analyses, which integrate along null geodesics back to the surface of last scattering, unless care is taken to ensure that the final result involves only first-order variables on the last scattering surface, which then only need be located to zero-order.  In this paper, we present a new expression for the CMB temperature anisotropy arising from linear scalar perturbations which is fully covariant and manifestly gauge-invariant. We obtain our expression by integrating the covariant and gauge-invariant Boltzmann equation~\\cite{LC-scalcmb,LC97-er} along observational null geodesics, paying careful attention to the gauge issues discussed above. Unlike some covariant results in the literature (see, for example,~\\cite{LC97-er,dunsby96b}), the expression derived here can be integrated trivially, in the instantaneous recombination approximation, up to a Rees-Sciama term in universes with arbitrary matter descriptions. (The covariant results in~\\cite{LC97-er,dunsby96b} can only be integrated in baryon-dominated universes, thus excluding CDM dominated universes, and other such models favoured by observation.) We base our treatment on the physically appealing covariant and gauge-invariant formulation of perturbation theory, as described in~\\cite{ellis89a,ellis89b}. In this approach, one works exclusively with gauge-invariant variables which are covariantly-defined and hence physically observable in principle. The covariant method has many advantages over other gauge-invariant approaches (such as that formulated by Bardeen~\\cite{bardeen80}). Most notably, the covariant variables have transparent physical definitions which ensures that predictions are always straightforward to interpret physically. Other advantages include the unified treatment of scalar, vector and tensor modes, a systematic linearisation procedure which can be extended to consider higher-order effects (the covariant variables are exactly gauge-invariant, independent of any perturbative expansion), and the ability to linearise about a variety of background models, such as Friedmann-Robertson-Walker (FRW) or Bianchi models.  For universes which are baryon-dominated at last scattering, our expression for the temperature anisotropy may be compared to other gauge-invariant analytic results in the literature. We show that, with suitable approximations, the result derived here reduces to that given by Panek~\\cite{panek86} and corrects a similar result given by Dunsby recently~\\cite{dunsby96b}. For the baryon-dominated universe, we use numerical results for the covariant, gauge-invariant variables on the last scattering surface, obtained from a gauge-invariant Boltzmann code~\\cite{LC-scalcmb}, to discuss the different physical contributions to the primary temperature anisotropy. In particular, we show that on intermediate and small scales, the ``monopole'' contribution to the temperature anisotropy is described by the spatial gradient of the photon energy density, in the energy-frame, on the last scattering surface. Since the (real) last scattering surface is approximately a surface of constant radiation temperature (so that recombination does occur there), the inhomogeneity of the radiation energy density in the energy-frame determines a distortion of the last scattering surface relative to the surfaces of simultaneity in the energy-frame. The extra redshift (due to the expansion of the universe) which the photons incur due to the distortion is seen as a ``monopole'' contribution to the temperature anisotropy on intermediate scales. There is a significant ``dipole'' contribution to the anisotropy on intermediate and small scales, which we discuss also.  We end with a discussion of the gauge issues inherent in the original Sachs-Wolfe calculation of the CMB anisotropy~\\cite{sachs67}, focusing on the ``monopole'' contribution to the temperature anisotropy on intermediate scales, described above. This contribution is often missed in Sachs-Wolfe type calculations through an incorrect treatment of gauge effects at the last scattering surface~\\cite{ellis-er97}. (Equivalently, the term is often missed through a failure to recognise the direction-dependence of the ``expected temperature'' used to define the temperature perturbation in many calculations.) This often neglected term, which is not important on large scales, is an essential component of the Doppler peaks in the CMB power spectrum.  We employ standard general relativity and use a $(+---)$ metric signature. Our conventions for the Riemann and Ricci tensors are fixed by $[\\nabla_{a},\\nabla_{b}] u^{c} = -{{\\mathcal{R}}_{abd}}^{c} u^{d}$, and ${\\mathcal{R}}_{ab} \\equiv {{\\mathcal{R}}_{abc}}^{b}$. We use units with $c=G=1$ throughout. ",
        "conclusions": "Starting from a covariant and gauge-invariant formulation of the Boltzmann equation, we have derived a new expression for the CMB temperature anisotropy under the instantaneous recombination approximation, valid for scalar perturbations in open, closed and flat universes. Our expression uses only covariantly-defined variables, and is manifestly gauge-invariant. The result is more useful in multicomponent models with scalar perturbations than earlier covariant results~\\cite{ellis-er97,LC97-er,dunsby96b}. In the case of a universe which is baryon-dominated at recombination, we find a simple expression for the anisotropy which corrects a similar result by Dunsby~\\cite{dunsby96b}. By making use of numerical solutions to the perturbation equations, we have discussed the conditions on the last scattering surface and their contributions to the characteristic features of the CMB power spectrum. We ended with a discussion of the original Sachs-Wolfe calculation for the temperature anisotropy. We have discussed why it is not necessary to locate accurately the last scattering surface in such calculations (because of the compensation effect), and how the extra term in the Sachs-Wolfe calculation, reported recently by Ellis and Dunsby~\\cite{ellis-er97}, is missed in many calculations which employ a gauge-dependent ``expected temperature'', since the angular dependence of this temperature is often overlooked. For a universe which is matter dominated at recombination, but not necessarily adiabatic, the extra term is the spatial gradient of the radiation energy density in the energy-frame, $(\\bar{{\\mathcal{X}}}_{k}^{(\\gamma)} Q^{(k)} /4)_{A}$. For models with isothermal surfaces of last scattering, this inhomogeneity describes a distortion of the last scattering surface relative to the surfaces of simultaneity of the energy-frame. The extra redshift incurred by this distortion is a significant component of the temperature anisotropy on intermediate and small scales."
    },
    "9804/astro-ph9804193_arXiv.txt": {
        "abstract": "We search OGLE-I photometric database for stars, which, as defined by formal criteria adopted by OGLE-I microlensing search, showed variability during only one out of 3 or 4 observing seasons. The results include 17 previously reported microlensing events, 2 newly discovered candidate events and 15 intrinsically variable stars that have a potential of contaminating samples of microlensing events. Based on photometry obtained in 1992 and 1993 OGLE \\#10 was tentatively included in the list of microlensing candidates, however its light curve in 1994 and 1995 shows many characteristics of the variable stars found in our search, and most likely it is not a microlensing event. For all stars which passed our tests, we provide $44 \\times 44$ arcsec ($101 \\times 101$ pixels) centered subframes from each OGLE-I frame in $I$ band. It is the first time when images used to derive photometry of microlensing events are available in convenient format to astronomical community. ",
        "introduction": "The search for rare cases of gravitational microlensing in the Local Group requires monitoring of $\\sim 10^6$ stars over several months in order to yield a significant rate of detections. A common implementation adopted is to construct a massive photometry database and subsequently select stars which experienced brightening of the type we expect on theoretical grounds (see Paczy\\'nski, 1996, for a recent review of basic theory, current microlensing searches and results). For the vast majority of events the light curves should follow a single point mass microlensing curve. Possible departures and exceptions from this basic case are extremely interesting. For random distribution of stars the probability that a given star will be lensed twice over duration of the experiment is negligible. Large fraction of binary lenses is expected to give raise to the population of ``wide binary'' light curves with two separate amplification regions ($\\sim 1\\%$ of the total rate of events). However in most cases the secondary peak should be weak ($A_{\\rm max} < 0.1$) and probably would only be discovered in light curves, which called attention because of the primary event (Di Stefano and Mao 1996). Photometry in very crowded fields, which are natural places to look for microlensing events, is often of limited quality. The majority of stars measured are just above the detection limit, many events have modest amplitudes, and numerous light curves are unevenly sampled. Therefore a microlensing curve may accidentally give a good fit to the brightening which is due to the intrinsic variability of the star. Certainly a repeated brightening of the same type would need to be very carefully examined before any claim of a detection of the wide binary microlensing repeater. As a result, it is generally hard to lower confusion rate without lowering the number of events returned by the procedure.  For microlensing events reported by OGLE-I project (Udalski et al. 1992) the basic requirement was that variability should occur during only one observing season (Udalski et al. 1994a). It was assumed that a sample of stars selected according to the above condition contains the majority of microlensing events and relatively few variable stars of other types. In this paper we investigate how many variable stars have light curves that, given the time sampling of OGLE-I experiment, appear constant during all seasons except for just one. The quantitative information about stellar variability background, against which microlensing events are detected, allows further tuning of the methods used in automated detection of the events. ",
        "conclusions": "A general conclusion is that variability background in OGLE-I search was reasonably well separated from microlensing events, although OGLE \\#10, a single candidate which most likely is not a microlensing event, constitutes a 5\\% confusion rate. A variable star, just like any other star, may be amplified by microlensing, but an increase of brightness by 0.1 mag may be naturally explained by intrinsic variability of this star, especially that similar objects (certainly not microlensing events) apparently exist.  We find two additional possible events that were not returned by the automated procedure of Udalski et al. (1994a). This is not surprising since we relax some of the selection cuts applied before. Moreover, one of those events happened near the end of the observing season while the other had very short time scale and poorly sampled light curve. Both of them are very inconspicuous.  The outburst experienced by MM3 $I$ 58214 (most likely a flare or CV star) is an important case which has a potential of contaminating samples of microlensing candidates. Suppose we had no data just before the flare. With photometric accuracy comparable to OGLE-I data such variable could be taken for a fading microlensing event and uncertainty would have to be resolved by spectroscopy and/or monitoring of the star long after the event.  We note a relatively large number of periodic or almost periodic variables which change amplitude. They mimic constant stars for time long enough to pass the most important criterion of the OGLE-I search, i.e., that a star should vary within a limited time interval with essentially constant flux at all other times. It is mostly due to relatively poor time coverage of the OGLE-I experiment and should not be difficult to overcome in the second phase of the project, OGLE-II (Udalski, Kubiak and Szyma\\'nski 1997). Two recommendations can be made for the future. First, low amplitude events, e.g. with $A_{\\rm max} < 1.5$, may be safely ignored in the calculation of the optical depth to prevent potential problems with objects similar to OGLE \\#10. MACHO team is already using such cut off to avoid contamination by ``bumpers''. Second, a requirement of roughly even photometric coverage of both sides of the magnification peak allows filtering out stars with (usually asymmetric) outbursts. OGLE-I events used in the optical depth determination satisfy such condition, however some of the events discovered by the EWS do not.  The full set of data used in this paper, including $101 \\times 101$ pix subframes extracted from every $I$ band image of each object in Tables~1 and 2, is available for public. Images in FITS format as well as photometric data in standard Johnson system may be retrieved via anonymous ftp from {\\tt astro.princeton.edu (128.112.24.45)} -- directory {\\tt /ogle/var\\_background} and {\\tt sirius.astrouw.edu.pl (148.81.8.1)} -- directory /ogle/var\\_background. See {\\tt README} file for details."
    },
    "9804/astro-ph9804299_arXiv.txt": {
        "abstract": "We present X-ray spectra obtained with \\sax\\ (Satellite per Astronomia X) of 10 BL Lac objects, selected from the Einstein Medium Sensitivity and Slew Surveys. We find that in about half of the  objects a fit in the 0.1-10 keV range with a single power law and free absorption yields values of $N_{\\rm H}$ larger than the Galactic ones. In most of these cases, however, broken power law fits with $N_{\\rm H}$ fixed at the Galactic values yield an alternative, better description of the data and indicate a steepening of the spectrum with increasing energy. One object (1ES1101-232) is detected up to $\\sim$ 100 keV. Its spectral energy distribution (SED) peaks in the medium energy X-ray band. For each object we compute the peak frequency of the SED from multifrequency data. The spectral indices $\\alpha_x$ in the 2-10 keV band ($F_\\nu \\propto \\nu^{-\\alpha_x}$) are smaller (i.e. flatter spectrum) for objects with higher peak frequencies. We therefore confirm and extend to higher energies the behavior already known for X-ray selected BL Lac objects in the ROSAT band. We do not find spectral indices smaller than 1; however, the flat distribution of $\\alpha_x$ and the correlation between $\\alpha_x$ and peak frequency found from our data suggest that a number of objects may exist, which in the quiescent status  have flatter spectrum and peak frequency in the hard X-ray range. ",
        "introduction": "\\noindent BL Lacertae objects are a rare type of Active Galactic Nuclei (AGN) characterized by strong and variable emission of non-thermal radiation across the entire electromagnetic spectrum, from radio waves to high energy $\\gamma$-rays. In three cases (Mkn 421: Punch et al. 1992; Mkn 501: Quinn et al. 1996; 1ES 2344+514: Catanese et al., 1998) the emission has been detected up to TeV energies. BL Lacertae objects comprise  the most violent (highly and rapidly variable, highly polarized) and most elusive (extremely difficult to find in optical surveys) sources amongst AGN. Unlike most other AGN they do not show evidence (by definition) for strong emission lines  or large Infra-Red or UV excesses. The emission from radio to $\\gamma$-rays can be explained as due to synchrotron radiation up to a certain maximum frequency (that ranges approximately from $10^{13}$ to  $10^{17}$ Hz), above which a sharp turnover occurs until a second component due to Compton scattered radiation dominates, making these objects detectable up to the highest energies so far accessible (see e.g., Ulrich, Maraschi and Urry, 1997). The extreme properties of BL Lacs require that the matter emitting the radiation moves at relativistic speeds in the direction of the observer.  The spectral change from synchrotron to Compton radiation is crucial for the understanding of the physics of BL Lacs. However, up to now this has been inferred only from the comparison of X-ray measurements carried out with different instruments and very often at different epochs. The wide energy band of \\sax\\ offers for the brightest objects the best opportunity to directly detect without ambiguity this spectral change and to study the X-ray spectra at the same epoch over a large interval.  To this end we have undertaken a program that aims at studying in detail the X-ray spectrum of a large and well defined subsample of soft X-ray selected BL Lacs. This sample includes mostly objects that are expected to show strong spectral curvature and spectral breaks, since the synchrotron break should occur just before or in the \\sax\\ band. We aim at measuring in detail the shape of the most energetic part of the synchrotron emission, and trying to establish where and how the Compton component becomes dominant. We also intend to look for the correlation between spectral slope and break energy found in ROSAT data (Padovani $\\&$ Giommi, 1996; Lamer, Brunner \\& Staubert, 1996). ",
        "conclusions": "We have analyzed the spectra for 10 X-ray selected BL Lacs observed with the Narrow Field Instruments on board the \\sax\\ satellite.  The sources are detected from $\\sim 0.2$ up to $\\sim$ 10 keV (and in one case up to $\\sim 100$ keV with the PDS instrument) with a very smooth appearance. The spectrum is generally well fitted by either a single power law, or by a broken convex power law that most probably represents the steepening after the synchrotron peak, whose position is determined also by using simultaneous optical observations.  Variability is not present during the short \\sax\\ exposure; analysis of ROSAT data shows for most of the sources little variability (within $\\sim 30\\%$) with respect to the \\sax\\ flux  and spectral indices consistent with the \\sax\\ ones. The spectral energy distributions, which include literature data, instead show variability in all bands.  The X-ray spectral indices $\\alpha_x$ range between 1 and 1.5 with a flat distribution and a mean value $\\langle \\alpha_x \\rangle = 1.31\\pm0.06$.  The scatter in the distribution is due to an anti-correlation we have found between $\\alpha_x$ and the frequency of the peak of the emission, $\\nu_{peak}$. This extends to the \\sax\\ band a correlation which had been discovered in the ROSAT band for this class of objects. The fact that sources with harder X-ray spectra have higher $\\nu_{peak}$ is expected if the \\sax\\ band is still dominated by synchrotron emission, which is also consistent with the spectral energy distributions of our BL Lacs.  Furthermore, we have no evidence of a spectral flattening (indicating the arising of the Compton component) in the present spectra, but future PDS detections, that are possible with exposure times slightly longer than those obtained here, might help in this respect.  The large fraction (at least 2 out of 10) of HBL selected in the soft X-ray band found with a flat ($\\alpha_x \\sim 1$) X-ray slope (i.e., they are near the peak of the synchrotron emission) and the distribution of $\\alpha_x$ values support the view that objects with even higher spectral peaks in their quiescent status indeed exist, and might be found in large numbers if we devise the correct strategy (e.g., samples at harder X-rays, TeV sources, etc.)  Moreover, these sources are good candidates to be TeV {\\it emitters}. In fact, in the sources with the flattest $\\alpha_x$ the peak of the synchrotron component is localized in the soft X--ray range. Electrons emitting at 1 keV by the synchrotron process have Lorentz factors $\\gamma\\sim 2.5\\times 10^5 (\\nu_{peak,1~keV}/B\\delta)^{1/2}$, where $\\nu_{peak} = 2.42 \\times 10^{17} \\nu_{peak,1~keV}$ Hz, $B$ is the value of the magnetic field in Gauss and $\\delta$ is the usual Doppler factor. Through the inverse Compton mechanism, they can emit up to $E\\sim \\gamma m_ec^2\\delta\\sim 130 (\\nu_{peak,1~keV}\\delta/B)^{1/2}$ GeV. If the magnetic and radiation energy densities are equal (as it is, approximately, in the three BL Lacs already detected in the TeV band), the flux level of the synchrotron and inverse Compton peaks is roughly equal. Low redshift sources are therefore good candidates to be {\\it detected} in the GeV--TeV band, while the high energy emission of the more distant sources could be absorbed in $\\gamma$--$\\gamma$ interactions with the background IR photon field, whose intensity is still uncertain. Indeed, a cutoff in the high energy spectrum could be used to determine the IR background (see, e.g. Stecker \\& De Jager, 1997).  This \\sax\\ project is still ongoing. We expect therefore to increase considerably the sample of soft X-ray selected BL Lacs for which we measure the spectrum in the 0.2-10 keV range and possibly above. With a larger complete sample, and combining the results with other complementary \\sax\\ projects, we expect to be able to draw a clearer picture of the relationship between the local X-ray slope and the overall energy distribution of this class of sources, in order to derive firmer conclusions on the behaviour at hard X-ray energies and on the mechanisms of the emission."
    },
    "9804/astro-ph9804066_arXiv.txt": {
        "abstract": "About 25\\% of the optical extragalactic sky is obscured by the dust and stars of our Milky Way. Dynamically important structures might still lie hidden in this zone. Various approaches are presently being employed to uncover the galaxy distribution in this Zone of Avoidance (ZOA). Results as well as the different limitations and selection effects from these multi-wavelengths explorations are being discussed. Galaxies within the innermost part of the Milky Way --- typically at a foreground obscuration in the blue of  $A_{\\rm B} \\ga 5^m$ and $|b| \\la \\pm5\\degr$ --- remain particularly difficult to uncover except for H\\,{\\sc i}-surveys: the Galaxy is fully transparent at the 21cm line and H\\,{\\sc i}-rich galaxies are easy to trace. We will report here on the first results from the systematic blind H\\,{\\sc i}-search ($v \\leq 12700$ km\\,s$^{-1}$) in the southern Zone of Avoidance which is currently being conducted with the Parkes Multibeam (MB) Receiver. ",
        "introduction": "To understand the dynamics within the local Universe -- the mass distribution and the local velocity field with its peculiar and streaming motions -- a detailed map of the 3-dimensional galaxy distribution is highly desirable. However, the dust extinction and confusion with stars in the disk of our Galaxy make this very difficult for $\\sim$25\\% of the sky, and the following questions remain unanswered:  Could a nearby Andromeda-like galaxy have escaped detection to date, hence change our understanding of the internal dynamics and mass derivations of the Local Group (LG), and the present density of the Universe from timing arguments (Peebles 1994)?  Is the dipole in the Cosmic Microwave Background Radiation (direction and amplitude) entirely explained by the gravity on the LG from the irregular mass/galaxy distribution? As the nearest galaxies ($v<300$ \\kms) generate 20\\% of the total dipole moment (Kraan-Korteweg 1989) nearby individual galaxies are equally important as massive groups, clusters and voids.  Is the mass overdensity in the Great Attractor (GA) region -- postulated from a large-scale systematic flow of galaxies towards ($\\ell,b,v)\\sim(320\\degr,0\\degr,4500$\\kms) (Kolatt \\etal\\ 1995) -- in the form of galaxies, hence does light trace mass?  Does the Supergalactic Plane, other superclusters, walls and voids connect across the Milky Way and might other large-scale structures (LSS) have gone undetected due to this 'zone of avoidance'? ",
        "conclusions": "The combination of the complementary multiwavelength surveys allow a new probing of LSS in the 'former' ZOA. The \\HI\\ surveys are particularly powerful at the lowest latitudes. But future merging of ZOA data with catalogs outside the ZOA will have to be done with care to obtain 'unbiased' whole-sky surveys.   From the sensitivity attained with the first 2 scans of the ZOA MB-survey it can be maintained that no Andromeda or other \\HI-rich Circinus galaxy is lurking undetected behind the extinction layer of the southern Milky Way.\\\\  {\\sl Acknowledgements} --- The help of the HIPASS ZOA team members R.D. Ekers, A.J. Green, R.F. Haynes, P.A. Henning, R.M. Price, E. Sadler, and L. Staveley-Smith is gratefully acknowledged."
    },
    "9804/astro-ph9804072_arXiv.txt": {
        "abstract": "Multicolour images of the starbursting metal poor blue compact galaxy ESO~338-IG04 have been obtained with the {\\it Wide Field Planetary Camera 2} on board the {\\it Hubble Space Telescope}. In the images we find numerous point-like sources concentrated towards the main body of the galaxy, which we identify as globular cluster candidates. We show that these objects are physically associated with the galaxy and that they are  spatially extended. Given their high intrinsic luminosities, these objects cannot be individual stars.  Using photometric evolution models we show that the objects constitute a rich population of  massive star clusters with ages ranging from a few Myr to $\\sim$ 10 Gyr, and masses ranging from  $10^4$ to more than $10^7 {\\cal M_{\\odot}}$. There are peaks in the age distribution of the clusters: one with objects $\\le30$ Myr, one at $\\sim 100$~ Myr, one at  $\\sim 600$~ Myr, one to two at $2.5-5$ Gyr and one at $\\sim10$ Gyr. The youngest objects are predominantly found in the crowded starburst region. They have  properties which agree with what is expected for young globular clusters, although it cannot be excluded that some of them may be dissolved or disrupted.  For objects older than a few times 10 Myr, the only plausible explanation is that these are globular clusters. The galaxy presently appears to be involved in a merger, which is the probable cause of the present globular cluster formation. The presence of a numerous intermediate age (2.5 to 5 Gyr) population of globular clusters, suggests that a previous merger might have occurred. As the starburst fades, this galaxy will become very rich in globular clusters. Transforming all objects to an age comparable to that of Milky Way globular clusters  reveals a luminosity function similar to the Galactic. We suggest that this galaxy is the result of a merger between a dwarf  elliptical and a gas rich dwarf. The possibility of dating the globular clusters offers an efficient way of studying the history of violent star formation in this and similar galaxies. ",
        "introduction": "Globular clusters (GCs) are generally old stellar systems and are believed to be the first objects to form in the process of the formation of a galaxy. This is supported by age estimates from observations of GCs in the Milky Way and other nearby galaxies. The Large Magellanic Cloud (LMC) is however known to host several young blue \"populous\" clusters, which could be young globular clusters. In the recent years blue globular cluster candidates have been found in some interacting/merging galaxies such as NGC~3597 (Lutz 1990, Holtzman et al. 1996), NGC~7252 (Whitmore et al. 1993), \"The Antennae\" (Whitmore \\& Schweizer 1995) and NGC~3921 (Schweizer et al. 1996). This has strengthened the idea  that GCs can form not only when galaxies form, but also when galaxies are reformed in the process of galaxy mergers (Schweizer 1986, Ashman \\& Zepf 1992, Whitmore 1996). This could possibly also circumvent the problem, faced by the idea that elliptical galaxies form by the merging of late type galaxies, that ellipticals have higher specific frequencies of globular clusters (van den Bergh 1984, 1994). Understanding how and when GCs can form will thus not only help us in understanding these systems, but also aid in understanding galaxy formation and evolution.    Studies of actively star-forming galaxies, e.g. Henize 2-10 (Conti et al. 1994), NGC~1244 (Hunter et al. 1994), NGC~1705 (O'Connell et al. 1994), NGC~1569 (De Marchi et al. 1997) and M82 (O'Connell 1995) have revealed the presence of star clusters with luminosities comparable to R126, the central cluster of 30 Doradus in the LMC, and higher. Meurer et al.(1995) used the HST for an ultraviolet study of nine starburst galaxies and  found bright \"super star clusters\" in all. These clusters are generally bluer than similar objects found in  interacting galaxies, and are preferentially found in the central regions of the starbursts. It has been proposed (e.g. Conti et al. 1994) that these blue objects might be forming globular clusters. This is however still an open question since it is not clear if these systems will survive as gravitationally bound systems. Uncertainties in the value of the (universal ?) stellar initial mass function (IMF) makes the mass estimates, mainly based on ultraviolet  data, highly uncertain, especially since some studies only includes one spectral bandpass. Many of these galaxies appear to be involved in some form of interaction.  One could interpret  these observations as evidence that massive star clusters can form in regions of very active star formation, such as giant extragalactic HII-regions (GEHRs)  like 30 Doradus. Kennicutt and Chu (1988)  made a statistical investigation of data on extragalactic blue populous clusters and giant HII-regions. They concluded that the young clusters in LMC are not luminous enough to evolve into globular clusters comparable to the massive galactic ones, and that far from all  GEHRs will produce young populous clusters that could become GCs.  Blue compact galaxies (BCGs) are characterised by their blue colours ($B-V \\le 0.5$), strong nebular emission lines, indicative of the formation of relatively hot (massive) stars, and low chemical abundances. The derived star formation rates could, considering the gas supply, only be sustained for a small fraction of a Hubble time. This together with the low metallicities (with IZw18 being the extreme in this sense)  once lead to the idea that BCGs might be genuinely young objects now experiencing their first star formation  epoch (Sargent and Searle 1970). Now, most BCGs are believed to be old, experiencing recurrent bursts of star formation intervened by long quiescent periods. Still we do not yet understand what triggers these bursts of star formation. One possible trigger mechanism is tidal or direct interactions with companion galaxies or gas clouds.  In this paper we will present conclusive evidence for young and old globular clusters in the blue compact galaxy ESO~338-IG04. Section 2 describes the observations and Sect. 3 describes how photometry was performed on the globular cluster candidates. In Sect. 4 we show that the objects are physically associated with the galaxy and that they are spatially resolved. In Sect. 5 we discuss how the observed photometric properties can be interpreted in terms of age and mass of the objects. Section 6 includes a further discussion on the nature of the cluster candidates, and Sect. 7 contains the conclusions.    \\subsection{General properties of the target galaxy}  ESO~338-IG04, also known as Tol~1924-416, resides at a distance of 37.5 Mpc ($v_{hel}=2813$ km~s$^{-1}$, $H_0= 75$ km s$^{-1}$ Mpc$^{-1}$; this value will be used throughout the rest of the paper) at the celestial coordinates $\\alpha_{1950}=19^{\\rm h}~24^{\\rm m}~30^{\\rm s}~ \\delta_{1950} = -41\\degr~34\\arcmin~00\\arcsec$. It  is intrinsically bright ($M_V=-19.3$) and blue ($B-V=0.4$), (Bergvall and \\\"Ostlin 1998). The oxygen abundance is $12 \\%$ of the solar value (Bergvall, 1985). It's physical size is small, $8.5 \\times 4.5$ kpc measured at $\\mu _V=25 {\\rm mag}~{\\rm arcsec}^{-2}$, the non dwarfish luminosity being due to the active star burst. It has an integrated HI-mass of   $3\\times10^9 {\\cal M_{\\odot}}$ (\\\"Ostlin et al. 1997b). The optical velocity field of ESO~338-IG04 show rotation aligned with it's apparent major axis, though with several large scale irregular features, and it is most easily understood as a merger between two galaxies or a galaxy and a gas cloud (\\\"Ostlin et al. 1997a). This is supported by the optical tail, extending towards the east. It has a spectroscopically confirmed companion galaxy (Bergvall unpublished; \\\"Ostlin et al. 1997a) which lies at a projected distance of 70 kpc, see Fig. 3. The companion is a somewhat fainter ($M_V = -17.9$), star forming galaxy which shows regular rotation (\\\"Ostlin et al. 1997a).  Bergvall noticed, in ground based images, an apparent concentration of faint blobs around the galaxy, which he interpreted as globular cluster candidates. We have re-observed  this galaxy with the HST and will in the subsequent sections show that there is substantial evidence that  this galaxy has formed massive globular clusters at several epochs.  \\begin{figure} \\picplace{8.5cm} \\caption[]{A Digitized Sky Survey (DSS) image (obtained with $Skyview$) showing the target galaxy and its companion. North is up and  east is left. The total angular size of the field is ~$9\\arcmin \\times 9\\arcmin$. ESO~338-IG04 is at the upper left, the companion at the lower right. } \\end{figure}  In July 1997, we obtained spectra of a few of the brightest globular cluster candidates, using the ESO New Technology Telescope (NTT), at La Silla, Chile. The results from this spectroscopic study will be presented in a future publication. ",
        "conclusions": "Multi-colour photometry with HST/WFPC2 of the metal poor blue compact galaxy ESO~338-IG04 (Tol~1924-416) has revealed a rich population of faint point-like sources in and surrounding the galaxy. Aperture photometry has been performed on these, and the photometric results have been transformed into ages and masses using a spectral evolutionary synthesis model. Special care was taken to assure that the sources are physically associated and not chance projected background or foreground sources. The results can be summarised as follows:  \\begin{enumerate}  \\item The objects discussed in this paper have absolute magnitudes, ~$M_v$, in the range -7 to -15 and $(v-i)$ colours ranging from less than -1 to almost 2.   \\item The vast majority of the found objects are spatially resolved star clusters, physically associated with the galaxy. The outer objects follow the luminosity distribution of the galaxy closely (Fig. 9).   \\item The number of interlopers, i.e. foreground stars, background galaxies and super giants in the target galaxy, in the presented sample is conservatively estimated to be less than ten objects. This leaves us with  more than 112 detected star clusters in  ESO~338-IG04.   \\item Using the photometric evolution model we show that the objects, in general, are well fitted by  Salpeter and Miller-Scalo IMFs, and that the resulting age distribution is insensitive to the adopted IMF. The ages of the clusters range from a few Myr to more than 10 Gyr.  \\item The objects have inferred masses  ranging from $10^4$ to more than $10^7 {\\cal M_{\\odot}}$.  \\item The above listed properties show that the objects are massive globular clusters of varying age.  \\item There are several peaks in the globular cluster age distribution. This shows that there have been several globular cluster forming events in this galaxy. These appear to have occurred ~$\\sim$ 11 Gyr ago, 2.5-5 Gyr ago, 600 Myr ago, 100 Myr ago and \"now\". The present event has lasted for a few times 10 Myrs.  \\item More than half of the total expected number of GCs have an age  in the range 2 to 5 Gyr. This peak consists of two subpopulations with either different age ($\\Delta_{\\rm AGE} \\sim 2.5$~Gyr) or metallicity ($\\Delta_{\\rm [Fe/H]} \\sim 1$~dex).   \\item The specific frequency, $S_N$, is high in this galaxy. Taking into account that the fading starburst will decrease the luminosity of the galaxy by 1.5 magnitudes in 1 Gyr, this galaxy will be come very rich in globular clusters, even if the newly formed GCs don't survive. The predicted specific frequency of globular clusters is much larger than for late type galaxies and comparable to that for giant ellipticals.    \\item We suggest that a merger is responsible for the current starburst and cluster formation. In view of the high specific frequency of old globular clusters suggests that an elliptical galaxy is a main ingredient in this system. To provide gas for star and cluster formation a gas rich dwarf must be the other main ingredient.  A major formation event 2.5-5 Gyr ago suggests that a merger involving one gas rich component also occurred at that time.  \\item In view of the high specific frequency and the present merger, we suggest that this galaxy will evolve into a moderately luminous elliptical galaxy, unless disturbed by possible future interactions.  \\item  This investigation for the first time gives strong support for a rich population both old and newly formed GCs in a metal poor blue compact galaxy.  \\item  Detection and dating of GCs in moderately distant BCGs is feasible using the HST and offers an efficient way of studying the violent star formation history, if also the effects of metallicity and extinction can be handled.  \\end{enumerate}"
    },
    "9804/astro-ph9804244_arXiv.txt": {
        "abstract": "Recent work suggests that rich clusters of galaxies commonly have large populations of dwarf (ie. low luminosity) members, that is their luminosity function (LF) turns up to a steep slope at the faint end. This population, or more particularly the relative numbers of dwarfs to giants, appears to be very similar for clusters of similar morphology, but may vary between cluster types. We have previously suggested that dwarfs may be more common in less compact, spiral rich clusters. Similarly we have found evidence for population gradients across clusters, in that the dwarf population appears more spatially extended. In the present paper we summarise the current evidence and propose, in analogy to the well-known morphology - density relation, that what we are seeing is a dwarf population - density relation; dwarfs are more common in lower density environments. Finally we discuss recent semi-analytic models of galaxy formation in the hierarchical clustering picture, which may give clues as to the origin of our proposed relation. ",
        "introduction": "Much recent work has been devoted to the question of the galaxy luminosity function (LF) within rich clusters, particularly with regard to the faint end which has become accessible to detailed study through various technical and observational improvements (see e.g., Driver et al. 1994; Biviano et al. 1995; Bernstein et al. 1995; Mohr et al. 1996; Wilson et al. 1997; Smith, Driver \\& Phillipps 1997 = Paper I; Trentham 1997a,b).  For the most part these studies concur that the LF becomes steep (Schechter (1976) slope $\\alpha \\leq - 1.5$) faintwards of about $M_{B} = -17.5$ or $M_{R} \\simeq -19$ (for $H_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$), and Paper I suggested that such a dwarf rich population might be ubiquitous. In a subsequent paper (Driver, Couch \\& Phillipps 1997a = Paper III) we have examined the luminosity distribution in and across a variety of clusters, examining the possible dependence of the dwarf population (in particular the ratio of dwarfs to giants) on cluster type and position within the cluster. In the present paper we summarise the evidence to date for the (dis)similarity of the dwarf population in different environments. ",
        "conclusions": "As with the corresponding morphology density relation for giant galaxies, the cause of our population - density relation could be either `nature' or `nurture', ie. initial conditions or evolution. Some clues may be provided by the most recent semi-analytic models of galaxy formation, which have been able to account in a general way for the excess of (giant) early type galaxies in dense environments (e.g., Baugh, Cole \\& Frenk 1996).  The steep faint end slope of the LF appears to be a generic result of hierarchical clustering models  \\footnote{ And was considered a problem until observational evidence for steep LFs increased!} (e.g., White \\& Frenk 1991; Frenk et al. 1996; Kauffmann, Nusser \\& Steinmetz 1997), so is naturally accounted for in the current generation of models. The general hierarchical formation picture envisages (mainly baryonic) galaxies forming at the cores of dark matter halos. The halos themselves merge according to the general Press-Schechter (1974) prescription to generate the present day halo mass function. However the galaxies can retain their individual identities within the growing dark halos, because of their much longer merging time scales. The accretion of small halos by a large one then results in the main galaxy (or cluster of galaxies, for very large mass halos) acquiring a number of smaller satellites (or the cluster gaining additional, less tightly bound, members).  Kauffmann et al. (1997) have presented a detailed study of the distribution of the luminosities of galaxies expected to be associated with a single halo of given mass. The LFs are somewhat disjoint owing to the specific halo masses modelled; especially for the low mass halo there is a preferred luminosity for the central galaxy plus a tail to lower luminosities. For a realistic mix of halo masses, these would no doubt be smoothed to look more like conventionally observed LFs. Nevertheless, we can still easily compare the numbers of dwarf galaxies per unit giant galaxy luminosity (rather than the amplitude of the giants' LF) between halos of different mass.  The Kauffmann et al. models mimic a \"Milky Way system\" (halo mass $5 \\times 10^{12} M_{\\odot}$), a sizeable group (halo mass $5 \\times 10^{13} M_{\\odot}$) and a cluster mass halo ($10^{15} M_{\\odot}$). Using their figure 2 (which also emphasises the identical faint end slopes predicted for all the different environments), we choose to quantify the number of dwarfs by $N_{-18}$, the number of dwarfs per system in the $M_{B} = -18$ bin. Because of the very similar slopes, the choice of bin or range of bins does not affect our conclusions, so this is the equivalent of the total number of dwarfs used in Figure 1. To quantify the giant population we choose the total light of galaxies of $M_{B} = -20$ or brighter, in units of $L_{*}$ galaxies (taking $M_{B}^{*} = -21$), which we call $N_{-21}$. Using this definition, rather than the actual number of galaxies brighter than some value (as in our observational data) allows for the discretization of the LFs for small halos. The results are summarized in Table 1. The ratio of these two values $N_{-18}$ and $N_{-21}$ then quantifies the relative dwarf galaxy populations. Roughly speaking, for smooth LFs with a shape similar to that observed, we should multiply these values by about 5, giving a range from about 1 to 3, to compare with our observational DGRs.  \\begin{table*} \\begin{center} \\begin{tabular} {lccccccccr} Halo Mass ($M_{\\odot}$) & $N_{-18}$ & $N_{-21}$ & $N_{-18}/N_{-21}$ \\\\ \\tableline $5 \\times 10^{12}$ & 0.2 & 0.34 & 0.6 \\\\ $5 \\times 10^{13}$ & 2.2 & 3.8 & 0.6 \\\\ $1 \\times 10^{15}$ & 40 & 190 & 0.2\\\\ \\end{tabular} \\end{center} \\caption{Dwarf numbers as a function of halo mass. \\label{tbl-1}} \\end{table*}   We see that the Milky Way and small group halos have similar numbers of dwarf galaxies per unit giant galaxy light, whereas the dense cluster environment has a much smaller number of dwarfs for a given total giant galaxy luminosity. Thus the predictions of the hierarchical models (which depend, of course, on the merger history of the galaxies) are in general agreement with our empirical results if we identify loose clusters and the outskirts of rich clusters with a population of (infalling?) groups (cf. Abraham et al. 1996), whereas the central dense regions of the clusters originate from already massive dark halos. By inputting realistic star formation laws etc., Kauffmann et al. can further identify the galaxies in the most massive halos with old elliptical galaxies, and those in low mass halos with galaxies with continued star formation. This would imply the likelihood that our dwarfs in low density regions may still be star forming, or at least have had star formation in the relatively recent past (cf. Phillipps \\& Driver 1995 and references therein). Note, too, that these galaxy formation models would also indicate that the usual (giant) morphology - density relation and our (dwarf) population - density relation arise in basically the same way. Finally, we can see that if these models are reasonably believable, then we need not expect the field to be even richer in dwarfs than loose clusters; the dwarf to giant ratio seems to level off at the densities reached in fairly large groups.  To summarise, then, we suggest that the current data on the relative numbers of dwarf galaxies in different clusters and groups can be understood in terms of a general dwarf population versus local galaxy density relation, similar to the well known morphology - density relation for giants. Low density environments are the preferred habitat of low luminosity galaxies; in dense regions they occur in similar numbers to giants, but at low densities dwarfs dominate numerically by a large factor. This fits in with the general idea that low luminosity galaxies are less clustered than high luminosity ones (particularly giant ellipticals). Plausible theoretical justifications for the population - density relation can be found within the context of current semi-analytic models of hierarchical structure formation."
    },
    "9804/gr-qc9804086_arXiv.txt": {
        "abstract": "Inflation of cosmic gauge and global strings is investigated by numerically solving the combined Einstein and field equations. Above some critical symmetry-breaking scales, the strings undergo inflation along the radial direction as well as the axial direction at the core. The nonsingular nature of the spacetimes around supercritical gauge and global strings is discussed and contrasted to the singular static solutions that have been discussed in the literature. ",
        "introduction": "Cosmic strings are linelike topological defects that may form as a result of a phase transition in the early universe. If a string is associated with a magnetic field, it is called a gauge string, otherwise it is a global string. They have attracted much attention because of their cosmological importance: deficit angle in the spacetime geometry and a candidate for the seed of structure formation in the early universe.  It was proposed that topological defects can inflate if the symmetry-breaking scale satisfies $\\eta \\gtrsim \\eta_c \\sim {\\cal O}(m_p)$ in Refs.~\\cite{Linde,Vilenkin}. This was later verified in numerical simulations by Sakai {\\it et al.}~\\cite{Sakai}. They found, in particular, that the critical value of $\\eta$ for domain walls and global monopoles is $\\eta_c \\simeq 0.33m_p$. Then what about cosmic strings? There is no reason why we exclude cosmic strings out of the topological inflationary category. Recently, it was numerically proved that a (2+1) dimensional gauge string can inflate by de Laix {\\it et al.}~\\cite{Tanmay}.  In this paper, we shall numerically solve the combined Einstein and field equations for a gauge and a global string in (3+1) spacetime dimensions. For the gauge string, we find that the core inflates if $\\eta \\gtrsim 0.25m_p$ with unit winding number in the critical coupling case (Bogomol'nyi limit). For the global string, $\\eta \\gtrsim 0.23m_p$. The critical values decrease as the winding number increases. For the gauge string, the critical value also decreases slightly as the coupling of the gauge field to the scalar field becomes weaker than the self coupling of the scalar field.  The asymptotic spacetime of a gauge string is known to be conical~\\cite{VilenkinC}. This spacetime exhibits a deficit angle $\\Delta =8\\pi G\\mu$, where $\\mu\\sim \\eta^2$ is the mass per unit length of the string. When the symmetry-breaking scale is sufficiently large, the deficit angle exceeds $2\\pi$ and analyses of the static solution show that the spacetime possesses a physical singularity outside the core of the string~\\cite{Gott,Ortiz,Laguna}. However, from our numerical results we know that supermassive strings are dynamical and undergo inflation at the core. Therefore, we believe that the static treatment of supermassive strings loses its validity and that we should treat them in a time-dependent way.  For global strings, the singularity exists regardless of the symmetry-breaking scale. Many people have tried to find a static solution of a global string and they found that there also exists a physical singularity outside the core of the string~\\cite{Cohen,Sikivie,Gregorys,Gibbons}. What was suggested to remove this singularity is again a time-dependent treatment of the string. Gregory~\\cite{Gregoryt} introduced a specific metric which has an axial time dependence  and showed that this spacetime is nonsingular.  In our work, we follow the evolution of supermassive gauge and global strings in a general time-dependent metric and show that no singularity develops in the spacetimes around the strings.  In the next section, we solve the Abelian Higgs model of a gauge string and discuss its inflation and spacetime geometry. Sec.~III is devoted to global strings. Our conclusions are summarized in Sec.~IV. In Appendix, we show the equations in detail and the numerical algorithms. ",
        "conclusions": "We have investigated inflation in cosmic strings. In the core region, the strings undergo inflation radially as well as axially when $\\eta \\gtrsim \\eta_c$. With unit winding number ($n=1$) the critical values for inflation were found to be $\\eta_c \\approx 0.25m_p$ for a gauge string in  the Bogomol'nyi limit ($\\beta=1$) and $\\eta_c =0.23m_p$ for a global string. The critical values decrease as $n$ and $\\beta$ increase. We have explained this $\\eta_c$ variation in terms of the core size of defects. The core of defects inflates when its size becomes bigger than the horizon scale: for larger $n$ and $\\beta$, strings have bigger cores, and the global string has a bigger core than the gauge string for a given $\\eta$. Regardless of the symmetry-breaking scale $\\eta$, around the center of defects the de Sitter expansion is established since the scalar field stays about the top of the potenital ($\\phi\\approx 0$). However, this is not sufficient for the cores of defects to inflate. Inflation requires another condition which is the core size being comparable to the horizon scale so that the core can be dynamical due to the gravitational effect. Or equivalently, the potential $V(\\phi)$ needs to be flat enough at $\\phi\\approx 0$ so that the field $\\phi$ can spend enough time about the top of the potential. For this condition to be satisfied, the symmetry-breaking scale $\\eta$ needs to be sufficiently large. This description also explains why we have somewhat lower critical values of $\\eta$ for strings than those for domain walls and global monopoles ($\\eta_c\\approx 0.33m_p$). Strings have bigger cores at the same symmetry-breaking scale than the other defects.  For supermassive gauge strings and all scale global strings, we have had troublesome physical singularities outside the core when we treat them in a static way. The elegant exit to nonsingular spacetimes is to introduce a time-dependent treatment. From the numerical simulations we could show that there is no singularity developing around time-dependent supermassive strings."
    },
    "9804/astro-ph9804008_arXiv.txt": {
        "abstract": "}[1]{{\\footnotesize \\noindent {\\bf Abstract} #1 \\\\}} \\renewcommand{\\author}[1]{\\subsubsection*{#1}} \\newcommand{\\address}[1]{\\subsubsection*{\\it#1}} \\setlength{\\textheight}{20cm} \\setlength{\\textwidth}{13.5cm}  \\begin{document}  \\chapter*{Galactic Environments of the Sun and Cool Stars\\footnote{To be published in {\\it Planetary Systems -- The Long View}, eds. L. M. Celnikier and J. Tran Than Van, Editions Frontieres, 1998}}  \\author{Priscilla C. Frisch} \\address{University of Chicago, Dept. Astronomy and Astrophysics, 5640 S. Ellis Ave., Chicago, IL  60637}  \\abstract{ The importance of understanding the current and historical galactic environments of cool stars is discussed.  The penetration of interstellar gas into a stellar astrosphere is a function of the interaction of the star with the interstellar cloud surrounding the star, and this factor needs to be understood if an efficient search for life-bearing planets is to be made.  For the Sun, both current and historical galactic conditions are such that if a solar wind were present, it would have excluded most inflowing interstellar matter from the inner regions of the heliosphere for the past few million years. Variations in heliosphere size over the recent historical path of the Sun are estimated, along with estimates of astrosphere sizes for selected nearby stars. Considering only possible effects due to encounters with interstellar clouds, stable planetary climates are more likely for inner than outer planets.} ",
        "introduction": "The Sun moves through space at a velocity of about 17 pc per million years.  This motion, combined with interstellar cloud motions driven by stellar evolution, yield a constantly changing galactic environment for the Sun and solar system.  This environment affects the interplanetary environments of both outer and inner planets in the solar system, including Sun--Earth coupling mechanisms. By analogy, the interactions between other cool stars and the galactic environment of that star needs to be understood as part of the process of identifying planets conducive to ``higher'' life forms.  The interstellar cloud surrounding the Sun at this time (known as the ``local interstellar cloud'', LIC), is warm, low density, and partially ionized: T$\\approx$7,000 K, n(H$^{\\rm \\circ})$$\\approx$0.2 cm$^{- 3}$, and n(e$^{-}$)$\\approx$0.1 cm$^{-3}$.  The standard assumption for diffuse interstellar clouds is that n(p$^{+}$)=n(e$^{-}$).  On the scale of typical cloud densities, the LIC is rather tenuous, and notably lower density than the 1 au solar wind density (see Fig. \\ref{densities}).  This accounts for the ability of the solar wind today to exclude most interstellar material from 1 au.  \\begin{figure}[ht] \\vspace*{4in} \\special{psfile=density.eps voffset=-107 vscale=60 hscale=60 } \\caption{Typical densities for material in our Galaxy. \\label{densities}} \\end{figure}  In this paper the basis for understanding the relation between the properties of stellar wind envelopes around cool star systems and the physical properties of the surrounding interstellar clouds are examined.  The author believes that the historical galactic environment of a star would have a direct impact on the stability of planetary atmospheres, and therefore on the distribution of intelligent life forms.  This conclusion rests partly on the observation that the solar system has been in a region of space virtually devoid of interstellar matter over the past several million years \\cite{fy86}, \\cite{journey}.  Additional reviews on interstellar matter (ISM) within the solar system can be found in the book {\\it The Heliosphere in the Local Interstellar Medium} \\cite{rudi}.  For more information on the properties of local ISM (LISM) see \\cite{fr95}.  For more information on nearby G-star space motions and environments, and the use of astrospheres as a test for interstellar pressure, see \\cite{fr93}. ",
        "conclusions": ""
    },
    "9804/astro-ph9804187_arXiv.txt": {
        "abstract": "We evaluate in a homogeneous way the optical masses of 170 nearby clusters ($z\\le 0.15$).  The sample includes both data from the literature and the new ENACS data (Katgert et al. 1996, 1998).  On the assumption that mass follows the galaxy distribution, we compute the masses of each cluster by applying the virial theorem to the member galaxies. We constrain the masses of very substructured clusters (about $10\\%$ of our clusters) between two limiting values.  After appropriate rescaling to the X-ray radii, we compare our optical mass estimates to those derived from X-ray analyses, which we compiled from the literature (for 66 clusters).  We find a good overall agreement.  This agreement is expected in the framework of two common assumptions: that mass follows the galaxy distribution, and that clusters are not far from a situation of dynamical equilibrium with both gas and galaxies reflecting the same underlying mass distribution.   We stress that our study strongly supports the reliability of present cluster mass estimates derived from X-ray analyses and/or (appropriate) optical analyses.   \\vspace*{6pt} \\noindent {\\em Subject headings: } galaxies: clusters: general - galaxies: distances and redshifts - X-rays: galaxies - cosmology: observations. ",
        "introduction": "The knowledge of the properties of galaxy clusters plays an important role in the study of large scale structure formation.  In particular, the observational distribution of the abundance of galaxy clusters as a function of their mass places a strong constraint on cosmological models (e.g., Bahcall \\& Cen 1993; Borgani et al. 1997; Gross et al. 1998; White, Efstathiou \\& Frenk 1993).  Moreover, recent studies stress the need for  having reliable estimates of cluster masses to constrain the ratio between the baryonic to total mass and the consequent value of $\\Omega_0$ (e.g., White \\& Frenk 1991; White et al. 1993b).  Indeed, the estimate of cluster masses is not an easy task, in spite of the various methods which are available.  The application of the virial theorem to positions and velocities of cluster member galaxies is the oldest method of cluster mass determination (e.g., Zwicky 1933).  More recent methods are based on the dynamical analysis of hot X-ray emitting gas (e.g., Cowie, Henriksen, \\& Mushotzky 1987; Eyles et al. 1991) and on gravitational lensing of background galaxies (e.g Grossman \\& Narayan 1989).   Mass estimates derived from the dynamical analysis of gas or member galaxies which are based on the Jeans equation or its derivations, such as the virial theorem, assume that clusters are systems in dynamical equilibrium (e.g., Binney \\& Tremaine 1987).  This assumption is not strictly valid; in fact, although clusters are bound galaxy systems, they have collapsed very recently or are just now collapsing, as is suggested by the frequent presence of substructures (e.g., West 1994).  However, some analyses suggest that the estimate of optical virial mass is robust against the presence of small substructures (Escalera et al. 1994; Girardi et al. 1997a; see also Bird 1995 for a partially different result), although it is affected by strong substructures (e.g., Pinkney et al. 1996).  Similar results come from studies based on numerical simulation (e.g., Schindler 1996a; Evrard, Metzler, \\& Navarro 1996; Roettiger, Burns, \\& Loken 1996) for X-ray masses estimated with the standard $\\beta$-model approach (Cavaliere \\& Fusco Femiano 1976), although some authors have claimed there is a systematic mass underestimation (e.g.; Bartelmann \\& Steinmetz 1996).  Dynamical analyses based on galaxies have the further drawback that the mass distribution or (alternatively) the velocity anisotropy of galaxy orbits should be known a priori.  Unfortunately, the two quantities cannot be disentangled in the analysis of the observed velocity dispersion profile, but only in the analysis of the whole velocity distribution which, however, requires a large number of galaxies (of the order of several hundreds; e.g. Dejonghe 1987; Merritt 1988; Merritt \\& Gebhardt 1994). Without some information from the relative distribution of dark and galaxy components, the virial theorem places only order-of-magnitude constraints on the total mass (e.g., Merritt 1987). The usual approach is to apply the virial theorem by assuming that mass is distributed like the observed galaxies (e.g., Giuricin, Mardirossian, and Mezzetti 1982; Biviano et al. 1993).  This assumption is supported by several  pieces of evidence coming both from optical (e.g., Carlberg, Yee, \\& Ellingson 1997) and X-ray data (e.g., Watt et al. 1992; Durret et al. 1994; Cirimele, Nesci, \\& Trevese 1997), as well as from gravitational lensing data, which, however, suggest a smaller core radius (e.g., Narayan \\& Bartelmann 1997).  The mass estimates derived from gravitational lensing phenomena are completely independent of the cluster dynamical status, but a good knowledge of cluster geometry is required in order to go from the projected mass to the cluster mass (e.g., Fort 1994).  Moreover, strong lensing observations give values for the mass contained within very small cluster regions ($\\lesssim$ one hundred of kpc) and weak lensing observations are generally more reliable in providing the shape of the internal mass distribution rather than the amount of mass (e.g., Squires \\& Kaiser 1996).  Up to now, few studies have dealt with wide comparisons between mass estimates obtained by different methods for the same cluster.  Wu \\& Fang (1996; 1997) found that masses derived from gravitational lensing analyses are higher than those from X-ray analyses by a factor of 2, but agree with those from galaxy analyses.  Indeed, mass estimates from lensing seem to agree with X-ray estimates when clusters are relaxed (e.g., Allen 1997).  However, Wu \\& Fang's works concern only clusters which lie at moderate redshifts and show gravitational lensing phenomena which could be enhanced in the presence of substructures (e.g., Miralda-Escud\\'e 1993; Bartelmann, Steinmetz, \\& Weiss 1995).  For nearby clusters, there is a trend to obtain larger masses from galaxy analyses than from X-ray analyses (e.g., Cowie et al. 1987; Mushotzky et al. 1995; David, Jones, \\& Forman 1995), but acceptable agreement exists in some individual cases (e.g., for the Coma cluster, Watt et al. 1992).  The classical approach of the virial theorem based on measurements of discrete velocities bears re-examining owing to the large new data sets which are now becoming available for nearby clusters, i.e. the ESO Nearby Abell Clusters Survey (ENACS) by Katgert et al. (1996, 1998). Moreover, the fair level of consistency among recent estimates of velocity dispersion of member galaxies resulting from different membership assignment procedures (cf. Fadda et al. 1996, hereafter F96, and Mazure et al. 1996) makes us confident of the robustness of our approach.  The aim of this work is to obtain reliable mass estimates.  These mass estimates will be used in the computation of the mass function of nearby clusters (Girardi et al. 1998).   The paper is organized in the following manner.  We describe the data sample and our selection procedure for cluster membership assignment in \\S~2. We briefly describe the methods used to compute cluster masses by using member galaxies in \\S~3. By assuming that mass follows the galaxy distribution, we compute virial mass estimates in \\S~4, and we verify their consistency with the results of the Jeans equation in \\S~5.  The strongly substructured clusters are analyzed in \\S~6.  We compare our mass estimates with those derived from X-ray analyses in \\S~7.  We discuss our results in \\S~8.  We give a brief summary of our main results and draw our conclusions in \\S~9.   Unless  otherwise specified,  we give errors at the 68\\% confidence level (hereafter c.l.)  A Hubble constant of 100 $h$ \\ks $Mpc^{-1}$ is used throughout. ",
        "conclusions": "The main points of this work may be summarized as follows:  i) We evaluate in a homogeneous way the optical masses of 170 nearby clusters ($z\\le 0.15$).  This sample, which is the largest set of clusters up to now analyzed in the literature, includes both data from the literature and the new ENACS data (Katgert et al. 1996, 1998).  ii) On the assumption that mass follows the galaxy distribution, we compute the masses of each cluster by applying the virial theorem to the member galaxies and we verify our results by using the Jeans equation.  iii) Our mass estimates are smaller than previous optical estimates. This fact is due both to our better membership assignment procedure and to the application of the correction due to the presence of the surface term in the virial theorem (recently stressed by Carlberg et al. 1997a).  iv) After appropriate rescaling to the X-ray radii, we compare our optical mass estimates to those derived from X-ray analyses, which we have compiled from the literature (for 66 clusters).  We find a good overall agreement.  v) The above agreement is expected on the basis of two common assumptions: a) that mass follows the galaxy distribution, b) that clusters are not far from a situation of dynamical equilibrium with both gas and galaxies reflecting the same underlying mass distribution. It should be pointed out that Carlberg et al. (1997a) have recently drawn similar conclusions for a sample of distant clusters (the CNOC sample).  In particular, we find evidence for a galaxy distribution which is colder and less extended than the gas distribution.  Several recent studies have casted doubts on cluster mass estimates and attempted to lower the cluster baryon fraction by reducing the cluster masses (e.g., Gunn \\& Thomas 1996; Wu \\& Fang 1996).  We stress that our study strongly supports the reliability of present cluster mass estimates derived from X-ray analyses and/or (appropriate) optical analyses.  Hence, it is even more difficult to reconcile present data with a $\\Omega_0=1$ Universe (e.g. White et al. 1993b).   Our cluster masses are suitable for statistical studies.  In particular, we did not reject a priori those clusters with a poor number of selected members, which usually have a small mass, in order to avoid having a final cluster sample biased towards more massive systems."
    },
    "9804/astro-ph9804323_arXiv.txt": {
        "abstract": "\\he\\ \\lya\\ \\lm 304/\\ha\\ \\lm 1640 emission lines are mainly produced by recombination, and their canonical ratio of $\\sim 10$ may be a sensitive reddening indicator. We obtain the high S/N optical spectra of two quasars and combine them with the far-UV spectra that show the \\he\\ \\lm 304 emission. For HS~1700+64, the \\he\\ \\lm 1640 emission is not detected, and an upper limit to it sets the ratio greater than 20. This may not be inconsistent with the theoretical value when all observational uncertainties are taken into consideration. For Q0302-003, the ratio is very low, on the order of unity. The most plausible cause for such a low ratio is extinction in the EUV band by very fine grains of dust. Q0302-003 has a prominent narrow component of $\\rm FWHM \\sim 2000 \\ \\kms$ in its major emission lines, and it appears that reddening is associated only with the line-emitting region. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part. ",
        "introduction": "The first UV spectroscopic observation of a quasar (3C~273, \\cite{davidsen77}) enabled a measurement of the \\lya/\\ha\\ ratio in an active galactic nucleus (AGN). The low ratio of $\\sim 1$ was in line with an independent study using the composite spectrum that was derived from various quasars (\\cite{baldwin}), but was in sharp conflict with a canonical ratio of $\\sim 10$ predicted by standard photoionization models (\\cite{deo}). This ``\\lya/\\ha\\ problem'' raised serious concerns about the validity of photoionization as the main line-emission mechanism in AGN and prompted extensive theoretical interest in the following years. Improved photoionization models (\\cite{ferland}, and references therein) invoke large column densities and moderate degrees of ionization. In a partially ionized zone, the low escape probability for \\lya\\ photons makes a high population of excited states, and collisional excitation from these levels enhances Balmer lines. Calculations using a reasonable AGN continuum lead to an enhancement of Balmer lines by a factor of $\\sim 2$, and hence may not fully explain the observed low \\lya/\\ha\\ ratio. Another explanation introduces intrinsic reddening in the line-emitting region (\\cite{dust}). The wavelength-dependent extinction reduces the intensities of observed UV lines, thus lowering the \\lya/\\ha\\ ratio. The observed Pa$\\alpha$/\\ha\\ ratio, however, appears to be too low for a straightforward full account as a reddening effect (\\cite{rick}). Significant evidence exists for dust in the narrow-line region of Seyfert galaxies (\\cite{deo}; \\cite{n93}) and for a decrease in extinction with increasing luminosity toward quasars (\\cite{cdz}; \\cite{rudy}). However, \\cite{wills} suggested that in intermediate-redshift quasars there may be significant reddening in the narrow-line region.  Accurate assessments of the reddening effect depend on the use of good line pairs whose intrinsic ratios are fairly stable. Since singly ionized helium is hydrogenic, the \\he\\ I(\\lya\\ \\lm 304)/I(\\ha\\ \\lm 1640) ratio should therefore be the same as that for hydrogen. The \\he\\ emission is produced mainly by recombination, because its excitation level of 40 eV is considerably higher than the average thermal energy in the line-emitting region. The wavelength of \\he\\ \\lm 304 emission coincides with that of O$^{++}$ transition $2p^2\\ ^3P_2 - 2p\\ 3d\\ ^3P_2^0$, allowing Bowen fluorescence radiation (\\cite{eastman}; \\cite{netzer}). This radiation mechanism, however, does not appreciably affect the \\he\\ ratio itself. The \\he\\ \\lm 304 emission should be extremely sensitive to reddening. While there are no hurdles in the theoretical aspects, it has taken some 20 years to advance from measuring this ratio in hydrogen to that in helium. ",
        "conclusions": "The \\ratio\\ ratio is quite different in these two quasars. Indeed the line profiles in Q0302-003 are narrower, making it easier to identify the weak \\he\\ \\lm 1640 feature. The S/N level is high enough that a \\he\\ \\lm 1640 feature should be detected even with a line width of $\\sim$ 12 000 \\kms. In Fig. 1 the profile of an assumed \\he\\ \\lm 1640 feature is plotted, with an intensity 20\\% of the \\he\\ \\lm 304, which should have been detected. It appears that difference is not simply attributable to line widths.  The intensity of \\he\\ \\lm 304 emission is affected by the Lyman line and continuum absorption by numerous intervening absorbers along the line of sight. This can be corrected if a high-resolution spectrum at longer wavelengths yields a list of absorption lines. Our estimate, based on the statistical result of \\cite{valley}, suggests an optical depth of 0.2 at a rest-frame wavelength of 300 \\AA\\ for a z=3.3 quasar. Therefore, this Lyman-Valley correction is not very significant.  The UV and optical observations are not simultaneous, and both quasars are probably variable. A comparison of the UV spectra of HS 1700+64 obtained between 1991 and 1995 finds a significant discrepancy in flux level, and that between the optical spectra taken between 1994 and 1996 shows that the \\lya\\ equivalent width varies by a factor of 2. The \\civ\\ equivalent width of Q0302-003 has increased by $\\sim 50\\%$ as compared with the data of Sargent, Steidel \\& Boksenberg (1989). Furthermore, the photometric quality of our optical spectra is questionable. A typical light loss with a small slit during an optical spectroscopic observation is $\\sim 15\\%$. These factors add uncertainties to the \\ratio\\ ratio. The derived line ratio is also subject to the reddening formulation. If we use the formula of Burstein \\& Heiles (1978), this ratio would be even lower.  We have carried out photoionization calculations (\\cite{cloudy}) with various parameters. With a broad range of the density, column density, flux and shape of the ionizing continuum,  the \\ratio\\ ratio varies within a narrow range between 9 and 11. It is therefore not practical to attribute the observed low value to special conditions in the line-emitting region. Note that a part of the \\he\\ \\lm 304 emission may receive a contribution  from \\oiii\\ \\lm 305 emission that is produced by Bowen fluorescence mechanism (\\cite{eastman}). If this were the case, the actual \\ratio\\ ratio would be even lower.  The low \\ratio\\ ratio in Q0302-003 may signal internal reddening in the line-emitting region. Dust grains with dimensions of $\\sim 3 \\times 10^{-6}$ cm are believed to produce Galactic extinction (\\cite{ext}) which generally follows a $1/\\lambda$ law. If intrinsic extinction is produced by even smaller grains, and the wavelength dependence of the extinction law applies to wavelengths as short as 300~\\AA, then an $E_{E - V} = 2.5$, where $E$ denotes a band around 300 \\AA, is needed to explain the discrepancy between the observed and theoretical ratio of \\ratio. This would translate into $E_{B-V} = 0.5$. If this is the case, significant presence of dust may be a reality in the broad-line region of some high-redshift quasars.  The quantitative formalism should be more complicated than that. While the extinction curves between 1 $\\mu$m and 1000~\\AA\\ can be approximated with a $1/\\lambda$ law, very little is known about the extinction properties below 1000~\\AA. \\cite{hawkins} and \\cite{martin} calculated the EUV extinction curve for graphite-silicate dust. Their results show {\\em decreasing} extinction from 1000~\\AA\\ to 100~\\AA. \\cite{pei} suggested that a numerical formula can be applied to other galaxies, possibly to those of higher redshifts, without assuming a Galactic dust-to-gas ratio. If such extinction is real, the same extinction affects the intensity of other UV lines as well. The intrinsic hydrogen \\lya/\\ha\\ ratio in these objects may actually be higher than the observed ones. Likely, the average ratio would be $ \\sim 6$, closer than the theoretical value of $\\sim 10$. This will, in turn, help understand the classical \\lya/\\ha\\ puzzle. If significant reddening does exist in the quasar broad-line region, the observed \\lya/\\ha\\ ratio may be corrected upward by an additional factor of $\\sim 2$.  Even for these two quasars, the I(\\lm 1216)/I(\\lm 304) ratio is very different. In HS~1700+64, this ratio is about 3, while in Q0302-003 it is about 30. Significant extinction in the EUV band can explain both the abnormal line ratios. The I(\\lm 1216)/I(\\lm 304) ratio in both objects is 30 or higher, consistent with photoionization models. Therefore, the likely cause for the low ratio in Q0302-003 is extinction in the broad line region by very small grains.  It may not be coincidental that narrow line widths and possibly significant reddening are present in the same object. Seyfert 2 galaxies often show a higher degree of extinction (\\cite{deo}), and the narrow-line region in Seyfert-1 galaxies generally exhibit a more significant reddening effect than the broad-line region (\\cite{luc}). The spectrum of Q0302-003 shows a significant narrow component with FWHM $\\sim 2000 $~\\kms\\ for major emission lines. Although this line width is not considered very narrow for Seyfert galaxies, it is for high-z quasars. Generally, narrow lines in high-z quasars (\\cite{sargent}) are not as common as in Seyfert galaxies. For example, \\cite{wills} found no detection of narrow line components in their radio-loud quasars of $\\rm 0.26 < z < 0.77$. They suggested a significant reddening with $E_{B-V} \\simeq 0.5$. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part.  Making an analogy of Seyfert galaxies and some low-redshift quasars, we suggest that Q0302-003 has a narrow-line region which contains a significant amount of dust. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part (\\cite{luc}).  Does reddening apply to the EUV continuum? Recent studies (\\cite{n95}; \\cite{bechtold}) found that the \\lya/\\hb\\ ratio ranges between about 1 and 40 and is approximately proportional to $f(1216)/f(4861)$, the ratio of continuum flux at adjacent points (\\cite{bechtold}). Such a correlation may suggest a possible reddening effect that applies to both the continuum and lines. In such cases, the equivalent widths of concerned lines should be fairly constant. Given the significant difference in the equivalent widths of \\he\\ lines in our quasar samples, we see no compelling reason that the continuum emission from the central source is heavily reddened. Significant reddening would also affect the intensities of infrared lines, and future studies of these lines may provide additional evidence for fine dust in the quasar environment."
    },
    "9804/astro-ph9804115_arXiv.txt": {
        "abstract": "The observed map of 1.809 MeV gamma-rays from radioactive $^{26}$Al (Oberlack et al, 1996) shows clear evidence of a Galactic plane origin with an uneven distribution.  We have simulated the map using a Monte Carlo technique together with simple assumptions about the spatial distributions and yields of $^{26}$Al sources (clustered core-collapse supernovae and Wolf Rayet stars; low- and high-mass AGB stars; and novae).  Although observed structures (e.g., tangents to spiral arms, bars, and known star-forming regions) are not included in the model, our simulated gamma-ray distribution bears resemblance to the observed distribution.  The major difference is that the model distribution has a strong smooth background along the Galactic plane from distant sources in the disk of the Galaxy.  We suggest that the smooth background is to be expected, and probably has been suppressed by background subtraction in the observed map.  We have also found an upper limit of $1 M_{\\sun}$ to the contribution of flux from low-yield, smoothly distributed sources (low-mass AGB stars and novae). ",
        "introduction": "The gamma-ray created by the decay of $^{26}$Al to $^{26}$Mg was the first discovered Galactic gamma-ray line (Mahoney et al 1982).  The $^{26}$Al nucleus decays by positron emission to the first excited state of $^{26}$Mg, which subsequently decays to the ground state emitting a 1.809~MeV gamma-ray.  The mean lifetime of $^{26}$Al, $\\tau = 1.05$ x $10^6$ years, makes the 1.809 MeV gamma ray line an excellent tracer for newly synthesized material released into the ISM over the last several million years.  The main production mechanism of $^{26}$Al is proton capture on $^{25}$Mg. Astrophysical environments that can produce $^{26}$Al include hydrostatic H-burning in the convective cores of massive stars and the H-burning shells of intermediate mass stars, and explosive H burning in novae.  The carbon and neon rich shells of massive stars are also a site for $^{26}$Al production both statically and explosively.  In addition to its production, the fresh $^{26}$Al must be transported into the ISM before it decays in order to be observable.  The explosive mechanisms present no problems, but the transport timescale in AGB stars is of similar order to the decay timescale causing a reduction in the amount of $^{26}$Al released into the ISM.  The first map of Galactic 1.809 MeV gamma-ray emission from $^{26}$Al was published by Oberlack et al (1996) from COMPTEL data.  This map has a 1$\\sigma$ angular resolution of $1.6^{\\circ}$, or $3.8^{\\circ}$ FWHM.  The map was produced using a Maximum-Entropy method after background subtraction. The 1.809 MeV gamma-ray map has several important characteristics, including the concentration of emission in the Galactic plane, a strong, irregular emission region toward the inner Galaxy, and a generally uneven, or clumpy, emission distribution. Along the Galactic plane there are several disconnected emission regions, some of which have been associated with O-B associations, spiral arm tangents, and the Vela SNR.  Chen et al (1996) also identify several of the regions with spiral arm tangents.  A recent and thorough review of the entire topic including observation, sources, and distribution of $^{26}$Al can be found in Prantzos and Diehl (1996).  These observations can best be explained with sources that are spatially concentrated and rare.  If the major sources of emission had small yields and a smooth Galactic distribution, the emission would be quite uniform. This is not seen in the published results (Oberlack et al 1996), which show large gaps along the Galactic plane between emission regions. We have therefore built a Monte Carlo model for the Galactic $^{26}$Al emission containing all potential astronomical sources.  We have also allowed the most massive stars to form clusters that do not dissociate in the lifetime of those stars. We have made only the simplest of assumptions about Galactic structure, an exponential disk and a bulge.  We have not attempted to represent any specific observed structures in the Galaxy.  All non-uniformities arise from the random nature of the simulation. This produces a map that, to the eye, has a strong resemblance to the observations, with the exception of a persistent uniform background not found in the reduced observational data.  We plan to use more statistically rigorous methods to measure the strength of this similarity in future work, as well as to consider some of the effects of non-uniform structure and the enhanced resolution of the INTEGRAL observatory. ",
        "conclusions": ""
    },
    "9804/astro-ph9804265_arXiv.txt": {
        "abstract": "The colossal power output of active galactic nuclei (AGN) is believed to be fueled by the accretion of matter onto a supermassive black hole. This central accreting region of AGN has hitherto been spatially unresolved and its structure therefore unknown. Here we propose that a previously reported `deep minimum' in the X-ray intensity of the AGN MCG$-$6$-$30$-$15, was due to a unique X-ray occultation event and that it probes structure of the central engine on scales $<  10^{14} \\ \\rm cm$, or $1.4\\times 10^{-7}$ arcseconds. This resolution is more than a factor of $\\sim 3\\times 10^{6}$ greater than is possible with current X-ray optics. The data are consistent with a bright central source surrounded by a less intense ring, which we identify with the inner edge of an accretion disk. These may be the first direct measurements of the spatial structure and geometry of the accreting black-hole system in an active galaxy. We estimate a mass lower limit for sub-Eddington accretion of $3.1\\times 10^{5} M_{\\odot}$. If the ring of X-ray emission is identified with the inner edge of an accretion disk, we get mass upper limits of $1.9\\times10^{8}$ and $9.1\\times10^{8} M_{\\odot}$ for a non-rotating and maximally rotating black hole respectively. We point out that our occultation interpretation is controversial in the sense that X-ray variability in AGNs is normally attributed to intrinsic physical changes in the X-ray emission region, such as disk or coronal instabilities. ",
        "introduction": "\\label{intro}  The accretion of matter onto a supermassive black hole as a mechanism for fueling the output of active galactic nuclei (AGN) is a paradigm strongly supported by recent spectroscopic observations of the iron K$\\alpha$ X-ray emission line (Tanaka \\etal 1995; Yaqoob \\etal 1995; Nandra \\etal 1997a and references therein). The extreme Doppler and gravitational energy shifts of the line photons, together with the shape of the line are consistent with an origin in a disk rotating about a black hole (Fabian \\etal 1989; Laor 1991). Strong gravitational redshifts, in which photon energies are changed by more than 10\\%, occur only when matter approaches closer than $\\sim  20$ gravitational radii ($= 20r_{g}; r_{g} \\equiv GM/c^{2}$) from a compact object. However, it is not possible to directly map the physical structure of the system since the highest spatial resolution of X-ray optics technology is a factor $\\sim 10^{6}$ too poor for even the closest AGN. Optical and radio observations provide greater resolution but the bulk of the emission at these wavelengths is not generated close enough to the central engine. So far, the highest spatial resolution observations, at radio wavelengths, have revealed a Keplerian disk in the AGN NGC 4258 down to only $\\sim 60,000 r_{g}$ (Miyoshi \\etal 1995; Maoz 1995). This still falls short by a factor $\\sim 3000$ of mapping the black-hole region.  The AGN MCG$-$6$-$30$-$15 ($z= 0.008$) was observed by the X-ray astronomy satellite {\\it ASCA} ({\\it Advanced Satellite for Astrophysics and Cosmology}; Tanaka, Inoue, and Holt 1994) for $\\sim 4.2$ days on 23 July 1994. Results from this observation have already appeared in the literature, including the 0.5--10 keV lightcurve (Iwasawa \\etal 1996, hereafter I96; Reynolds 1997; Yaqoob \\etal 1997; see Figure 1) and a broad, asymmetric, variable iron K line with a strong red wing,  consistent with a disk inclined at $\\sim 30^{\\circ}$ rotating about a black hole (Tanaka \\etal 1995; I96). The X-ray luminosity exhibits erratic variability on all timescales down to $<50$ s (Matsuoka \\etal 1990; Green, McHardy, and Lehto 1995; Reynolds \\etal 1995; Nandra \\etal 1997b). Causality arguments alone cannot put constraints on the size of the X-ray emission region since the high-frequency, lower amplitude variability may occur at localized regions of the source. Figure 1 shows an extended intensity dip at the end of the observation, from $\\sim 3.3\\times10^{5}$s to $\\sim 3.6\\times10^{5}$s. This feature has previously been dubbed as the 'deep minimum', or DM (I96). A closer inspection (Figure 2a) reveals a remarkable (albeit approximate) symmetry about the minimum luminosity.  We propose that the dip was caused by an occultation of the X-ray source by optically-thick matter. This interpretation is controversial and 'non-standard', as X-ray variability in AGNs is normally attributed to intrinsic properties of the X-ray emission region, such as disk or coronal instabilities. However, so little is know about structure of the central engine in AGNs that the occultation scenario should be explored. The obscurer must be optically thick because the dip continuum spectrum only shows evidence for {\\it weak} absorption, nowhere near enough to explain the observed intensity variation over the whole \\asca bandpass (see Weaver \\& Yaqoob 1998, hereafter WY98). The luminosity at the absolute minimum of the dip is $\\sim 0.4$ of the pre-dip value and must represent persistent emission which has much smaller surface brightness than the primary source. Hereafter we will refer only to the primary X-ray emission, unless explicitly referring to the persistent emission. The proposed obscurer very likely hides the most compact and variable part of the X-ray source, since the usual rapid variability outside the dip is absent during the obscuration.  Of course, it is possible that the dip is due to intrinsic variation of the X-ray source. However, the origin of X-ray variability in AGN is not understood. Models which come close to successfully reproducing the observable quantities obtained from AGN lightcurves are of the shot-noise or `rotating hot-spot' variety (e.g. Green \\etal 1993; Bao and Abramowicz 1996 and references therein). However, the parameters of such models must be highly tuned in order to reproduce AGN power spectra. On the other hand, we show in this paper that a very simple-minded model can account for the temporal profile of the intensity dip in MCG $-$6$-$30$-$15 and briefly discuss the implications for AGN X-ray variability in general. ",
        "conclusions": "If accreting matter is exposed to the same UV/X-ray luminosity that we observe ($L \\sim 4 \\times 10^{43} \\ \\rm ergs \\ s^{-1}$) then for gravitational infall to overcome outward radiation pressure requires the mass of the central black hole, $M_{\\rm BH}$, to exceed $3 \\times 10^{5} M_{\\odot}$. Thus, if $r_{1}$ is the radius of the inner edge of the accretion disk, identified with the last stable orbit of matter, and $\\Delta t$ is the time taken for the obscurer to traverse $r_{1}$ (i.e. $t_{3}-0.5[t_{2}+t_{1}$]), then $\\kappa r_{g}=r_{1}$ where $\\kappa=6$ or 1.24 for Schwarzschild or maximally rotating Kerr metrics respectively. But, $ r_{1}<c \\Delta t$ and $r_{g} = 1.48 \\times 10^{13} \\ \\rm (M_{\\rm BH}/M_{\\odot})$ cm, so $M_{\\rm BH} < 1.9 \\times 10^{8} M_{\\odot}$, or $<9.1\\times 10^{8} M_{\\odot}$, for a Schwarzschild or Kerr black hole respectively. Assuming instead, the Keplerian velocity at the inner disk edge as the maximum velocity of the obscurer ($v/c = \\sqrt{r_{g}/r} = \\sqrt{1/\\kappa}$) yields smaller mass upper limits of $M_{\\rm BH}<7.7\\times10^{7} M_{\\odot}$ and $M_{\\rm BH} <8.2\\times10^{8} M_{\\odot}$ for a Schwarzschild and extremal Kerr metric respectively. If the obscurer is at a distance, $d$, from the central source, $r_{1} = v \\Delta t = c \\sqrt{r_{g}/d}$ gives $d = c^{2}[\\Delta t]^{2}/\\kappa^{2}r_{g} = 6 \\times 10^{15}[\\Delta t]^{2}/[\\kappa^{2} (M_{\\rm BH}/M_{\\odot})]$ cm. Thus, $M_{\\rm BH} > 3 \\times 10^{5} M_{\\odot}$ implies $d<1.8\\times 10^{16}\\ \\rm cm$ and $d<4.2\\times10^{17}\\ \\rm cm$ for Schwarzschild and Kerr metrics respectively. The origin of the optically-thick blobs is unspecified but Guilbert and Rees (1988) presented some simple arguments for the existence of dense ($n > 10^{15} \\rm \\ cm^{-3}$), optically thick matter residing at the heart of accreting sources. The only independent estimate of the size of the `blobs' is that they should be much thicker than $10^{9}/n_{15} \\ \\rm cm$ ($n_{15}$ in units of $10^{15} \\rm \\ cm^{-3}$). The blobs must be optically thick even near their physical boundaries (i.e they must have fairly sharp edges), otherwise the dip profile would not be so well defined. Also, the blobs must be fairly stable, esepcially if they are created in the central region itself and 'propelled' up to high altitudes.  The nature of the bright central X-ray source is intriguing. An inclined jet is unlikely, since even at $30^{\\circ}$ the inflexions in the dip profile ($t_{2}$--$t_{3}$ and $t_{6}$-$t_{7}$) would have different durations. If the central X-ray source extracts its energy directly from the black hole then the metric is likely to be Kerr since energy cannot be extracted from a non-rotating black hole (Blandford and Znajek 1977; see also Ghosh and Abramowicz 1997). Occultations such as the one described here may occur frequently in AGN, but the relative sizes of the obscurer and source must be just right in order to observe such a clear event. Indeed the usual rapid variability or flicker, may be partly caused by the transit of optically thick bodies smaller than the source. We tested this hypothesis, again using a simple-minded model. Representing the bright central source as a circular disk with uniform emissivity (ignoring the ring due to its weaker emission), optically-thick blobs with ranges in radii (relative to the source) and velocities taken from Gaussian distributions were passed over the source. An additional parameter is required to specify the 'blob birth-rate' (i.e. the rate at which new blob trajectories are started). Such a model was used to produce predicted lightcurves for different model parameters. The power spectrum of each lightcurve was computed using the method of Papadakis and Lawerence (1993), omitting Poisson noise. In the range $\\sim 10^{-5}$ Hz to $\\sim 10^{-2}$ Hz, no preferred or 'universal' power-law spectral slope was found. It is possible to produce power-law spectra with slopes similar to those typically measured ($\\sim -1$ to $-2$) but for most parameter values, the slopes are too steep. Thus, fine-tuning would be necessary to explain the `universal' power-law slopes found in the handful of AGN in which it can be measured (e.g. Lawerence and Papadakis 1993; Green \\etal 1993). This is essentially because if the product of blob crossing-time and birth-rate is too large or too small, there will be no variability. The direct simulated lightcurves assume, of course that the {\\it intrinsic} source intensity is constant, which almost certainly is not the case. Thus our model does not explain AGN variability in general but its effects are potentially important to consider in any model of AGN variability.  We thank the \\asca TEAM and mission operations at ISAS, Japan, for their efforts and hard work; Kim Weaver for her work on the dip spectrum, and Karen Leighly, Paul Nandra for some useful discussions. We also thank the anonymous referee. This research made use of archival data at the HEASARC, Laboratory for High Energy Astrophysics, NASA/Goddard Space Flight Center."
    },
    "9804/astro-ph9804053_arXiv.txt": {
        "abstract": "We investigate different types of neutrino hot dark matter with respect to structure formation and anisotropies in the cosmic microwave background radiation (CMBR). The possibility of neutrino hot dark matter produced through the decay of a heavier neutrino by the process $\\nu_H \\to \\nu_L + \\phi$, where $\\phi$ is a scalar particle, is discussed in detail. This type of dark matter can possibly be distinguished observationally from the standard neutrino dark matter by using new CMBR data from the upcoming satellite missions MAP and PLANCK. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804029_arXiv.txt": {
        "abstract": "The importance of the interstellar magnetic field is studied in relation to the evolutions of superbubbles with a three-dimensional (3D) numerical magnetohydrodynamical (MHD) simulation. A superbubble is a large supernova remnant driven by sequential supernova explosions in an OB association. Its evolution is affected by the density stratification in the galactic disk. After the size reaches 2--3 times the density scale-height, the superbubble expands preferentially in the $z$-direction. Finally it can punch out the gas disk (blow-out). On the other hand, the magnetic field running parallel to the galactic disk has an effect to prevent from expanding in the direction perpendicular to the field. The density stratification and the magnetic fields have completely opposite effects on the evolution of the superbubble. We present results of 3D MHD simulation in which both effects are included. As a result, it is concluded that when the magnetic field has a much larger scale-height than the density, even for a model that the bubble would blow out from the disk if the magnetic field were absent, the magnetic field with the strength of 5 $\\mu$G can confine the bubble in $|z| \\alt 300$ pc for $\\simeq$20 Myr (confinement). In a model that the field strength decreases in the halo in proportion to $B \\propto \\rho^{1/2}$, the superbubble eventually blows out like a model of $B=0$ even if the magnetic field in the mid-plane is as strong as $B=5\\mu$G. ",
        "introduction": "A superbubble is a complex consisting of an OB association, surrounding X-ray emitting hot gas, and a corresponding HI hole/shell. Three examples are well-known in our Galaxy: Cygnus (Cash et al. 1980), Orion-Eridanus (Cowie,  Songaila \\& York 1979; Reynolds \\& Ogden 1979), and Gum nebula (Reynolds 1976). Superbubbles are found as HI shells and holes in external galaxies, such as LMC (Meaburn 1980; Dopita, Mathewson \\& Ford 1985), M31 (Brinks \\& Bajaja 1986), M33 (Deul \\& den Hartog 1990), M101 (Kamphuis, Sancisi \\& van der Hulst 1991), and so on. Since the sizes of these objects are in the range of 100 pc -- 1kpc, this can not be explained by a single supernova explosion [the size of an ordinary supernova remnant (SNR) is $\\alt 50$pc]. The amount of energy required for such a superbubble reaches $5\\times 10^{51} {\\rm erg} - 10^{54} {\\rm erg}$ (Tenorio-Tagle \\& Bodenheimer 1988). There are two models proposed for formation of superbubbles: (1) a large SNR driven by sequential supernova explosions in an OB association and (2) a complex formed by a collision of a high-velocity cloud and the galactic disk (Tenorio-Tagle 1991). Here, we confine ourselves to the first model and discuss the evolutions. For review papers of this field, see Tenorio-Tagle \\& Bohdenheimer (1988), Spitzer (1990), Tomisaka (1991), Bisnovatyi-Kogan \\& Silich (1995).  Here, we summarize the evolution very briefly. After the size of bubble exceeds the density scale-height, the bubble becomes elongated in the direction perpendicular to the galactic disk (Tomisaka \\& Ikeuchi 1986; Tenorio-Tagle, Bodenheimer \\& R\\'{o}\\.{z}yczka, 1987; MacLow \\& McCray 1987). It is shown that when the mechanical luminosity released by sequential supernova explosions is high, the expansion of the bubble to the halo is accelerated and the hot gas contained in it flows into the galactic halo (galactic fountain), finally. However, the magnetic fields running parallel to the galactic disk prevents the gas from flowing in the vertical direction (perpendicular to the fields).  Tomisaka (1990) studied the adiabatic evolution of a superbubble in uniform magnetic fields and showed that the bubble driven by a mechanical luminosity of $L_{\\rm SN}\\simeq 3\\times 10^{37} {\\rm erg\\,s^{-1}}$ is confined in the galactic disk by the effect of the magnetic field provided its strength is as large as $B_0=5\\mu$G. This shows the superbubble is confined in the galactic disk, if (1) the magnetic fields have a large scale-height, $H_B$, and (2)  $L_{\\rm SN}\\simeq 3\\times 10^{37}{\\rm erg s^{-1}}$ and  $B_0\\agt 5\\mu$G. However, this result may be affected by the assumption of the adiabatic gas. The interstellar magnetic fields seems to play an important role in the evolution of a superbubble (see also Mineshige, Shibata, \\& Shapiro 1993, Ferriere, MacLow, \\& Zweibel 1991).  In the present paper, a full magnetohydrodynamical calculation has been done including the radiative cooling. Further, the effect of  distributions of magnetic field strength is studied. Plan of the present paper is as follows: in section 2 is given an analytical estimate to a threshold  mechanical luminosity under which the superbubble is confined in the galactic disk. From this estimation, we choose the values for the parameters $L_{\\rm SN}$, $B_0$, and $H_B$. Numerical method and other assumptions are also presented in $\\S$2. Section 3 is for the numerical result, in which the effect of the magnetic fields is shown. In section 4,  we will discuss the observability of the superbubble in external galaxies. Using the evolution obtained, we will show that a large fraction of the interstellar space is occupied with superbubbles. ",
        "conclusions": "\\subsection{Blow-out or confinement?}  As shown in the preceding section, although in Model B the mechanical luminosity is much larger than the critical luminosity of equation(\\ref{eqn:Lcr0}), the vertical expansion is much decelerated by the effect of magnetic fields. This shows that $L_{\\rm SN} < L_{\\rm crit}$ may be only a sufficient condition for confinement of the superbubble.  The expansion of a spherical shell driven by a pressure in the hot cavity $p$ is formulated as follows: \\begin{equation} \\frac{dMv}{dt}=4\\pi\\, R^2 (p-p_{\\rm out}), \\end{equation} \\begin{equation} \\frac{dM}{dt}=4\\pi\\, R^2\\, v\\, \\rho_{\\rm out} \\end{equation} \\begin{equation} \\frac{dR}{dt}=v, \\end{equation} \\begin{equation} \\frac{dE}{dt}=L_{\\rm SN}-4\\pi R^2\\, p\\, v, \\end{equation} where, $M$, $R$, $v$, $p_{\\rm out}$, $\\rho_{\\rm out}$ and $L_{\\rm SN}$ represent, respectively, the mass of the shell, the radius and velocity of the shell, the interstellar pressure and its density and the energy release rate from an OB association. These equations are, respectively, the equation of motion, the mass conservation, the relation of a size $R$ to a velocity $v$, and the first law of thermal physics. If we assume $p=(2/3)\\times E /(4\\pi R^3/3)$, these four equations can be solved numerically. Figure 9 shows a resultant expansion law of a spherical superbubble with $L_{\\rm SN}=3\\times 10^{37}{\\rm erg~s^{-1}}$ in a {\\em uniform interstellar medium} of $n_0=0.3{\\rm cm^{-3}}$. Each curve corresponds to different external pressures as $p_{\\rm out}=1.7\\times 10^{-12}{\\rm erg~cm^{-3}}$ (solid line), $p_{\\rm out}=1\\times 10^{-12}{\\rm erg~cm^{-3}}$ (dotted line), and $p_{\\rm out}=0$ (dashed line). Weaver et al.'s (1977) solution, equation (1), agrees with the curve of $p_{\\rm out}=0$. This figure shows that  the interstellar pressure plays an important role especially in the late phase of the evolution. This is understood as follows: in the late phase of the superbubble the difference between the internal pressure and the outer one is small and this small difference drives the shell further. Therefore, the critical luminosity would be  underestimated if we use a solution without taking  the outer pressure into account. The shell of a superbubble continues to expand as long as the energy ejection continues, while a supernova remnant stops its expansion after $p_{\\rm out}=p$. Thus, exactly speaking, the superbubble is never confined as long as the OB association is alive. However, a slow expansion driven by a small pressure difference as $p-p_{\\rm out}$ is considered as a signature of the confinement.  In this figure, we also plot an equivalent radius, which is defined using the volume occupied with a hot matter ($V_{\\rm hot}$) as \\begin{equation} R_{\\rm equiv}\\equiv \\left(\\frac{3V_{\\rm hot}}{4\\pi}\\right)^{1/3}, \\label{Requiv} \\end{equation} for Models A and C. The equivalent radius for Model A in which the bubble is almost confined in the galactic gaseous disk shows a similar expansion law to that obtained by a thin-shell model (a solid curve). In contrast, that of Model C indicates a completely different expansion law such that after $t \\agt 10$ Myr the equivalent radius increases rapidly and in $t\\simeq 35$Myr $R_{\\rm equiv}$ surpasses the model with $p_{\\rm out}=0$. These differences seem to come from the distribution of magnetic field strength. In Model A the total pressure (thermal plus magnetic one) is almost constant as the bubble expands, because the magnetic pressure is dominant over the thermal one and magnetic fields are uniform. While, in Model C the total pressure drops according to the density distribution $\\rho(z)$. This figure indicates that when the equivalent radius is well fitted by this thin shell model the bubble is nearly confined in the disk even if the gas disk has a finite scale-height. In contrast, if hot gas is ejected from the galactic disk, the equivalent radius shows a more rapid expansion than that derived by this thin-shell model. Values of equivalent radii at the epochs when numerical runs end are shown in table 1.  \\subsection{Observability}  A shear motion in the galactic rotation and rotation itself may play a role in the evolution of a superbubble (Tenorio-Tagle \\& Palou\\v{s} 1987; Palou\\v{s} et al. 1990;  Silich 1993). Galactic shear seems to deform the shape and the Coriolis force makes the shell rotate. The characteristic time-scales of the rotation, $\\tau_R$, and the shear, $\\tau_S$ are estimated respectively as \\begin{equation} \\tau_{\\rm R} \\sim 1/\\Omega_0 \\sim 40{\\rm Myr} (\\Omega_0/26{\\rm km~s^{-1}~kpc^{-1}})^{-1} \\end{equation} and \\begin{equation} \\tau_{\\rm S}\\sim (l d\\Omega/dR)^{-1} \\sim 320 {\\rm Myr} (l/1{\\rm kpc})^{-1}(\\Omega_0/ 26{\\rm km~s^{-1}~kpc^{-1}})^{-1} (R_0/8.5{\\rm kpc}), \\end{equation} where $\\Omega_0$, $R_0$, and $l$ are the angular speed of galactic rotation, distance from the galactic center, and a typical size of a superbubble, respectively. Since active  SN-explosion phase continues for $\\sim 50$ Myr for an OB association (McCray \\& Kafatos 1987), in the late phase of $\\tau_{\\rm R}\\la t \\la 50$Myr, the effect of the Coriolis force seems to appear as a deformation force of the shell. The $\\alpha\\omega$-dynamo mechanism driven by superbubble was studied recently by Ferri\\`{e}re (1992).  The  galactic rotation has a little effect on the evolution of a superbubble. Thus, if shells or holes observed in external galaxies are elongated after their inclinations are corrected, their direction seems to indicate that of the magnetic fields. There have been listed 141 HI holes in M31 by Brinks \\& Bajaja (1986). These holes are observed, more or less, as elliptical. Particularly, the holes found near the major axis of M31 are important to determine the physical shape of the holes. Many of these HI holes  have their major axes perpendicular to the galaxy's major axis. Since this is not explained by projection due to the inclination of M31 ($i=77\\deg$), these seem to have physically a shape elongated along the azimuthal direction of the galaxy. This is not inconsistent with  observations indicating that a global pattern of the magnetic field is ring-like in M31, which is measured by  radio  linear polarization observations (for a review, see Sofue, Fujimoto \\& Wielebinski 1986), in other words, magnetic field lines run in the azimuth direction in M31.   \\subsection{Porosity}  If  hot gas contained in superbubbles occupies a large volume of the galactic disk, a picture of the interstellar medium should be changed (McKee \\& Ostriker 1977). The fraction of areas covered by superbubbles younger than $\\tau_{\\rm active}$ is estimated with a quantity called as two-dimensional porosity which is defined as \\begin{equation} Q(t < \\tau_{\\rm active}) \\equiv r_{\\rm OB}\\int_0^{\\tau_{\\rm active}} S(t)dt, \\end{equation} where $r_{\\rm OB}$ is the formation rate of OB associations per unit area and $S(t)$ represents the area which covered by a hot cavity on the mid-plane of the disk $z=0$. This is identical with a two-dimensional porosity parameter calculated by  Heiles (1990). He estimated galaxy-wide average of two-dimensional porosity $Q_{\\rm 2D}\\simeq 0.30$. $\\tau_{\\rm active}$ should be chosen equal to the oldest age of a superbubble which contains a hot gas inside. If we assume $\\tau_{\\rm active}=20$ Myr and integrate $S(t)$ for Models A and C ($S\\equiv \\pi x_c y_c$),  these two models give respectively $2.08\\times 10^6\\, {\\rm pc^2~Myr}$ and $1.57\\times 10^6\\,{\\rm pc^2~Myr}$. We adopt the estimation of $r_{\\rm OB}$ from a galactic type II supernova rate of $r_{\\rm II}\\sim 0.01{\\rm yr^{-1}}$, that is, we assume that all type II SNe occur in OB associations, number of type II SNe in an OB association is constant irrespective of richness of association as $N_{\\rm SN}\\sim 100$ and  OB associations  are uniformly distributed in the galactic disk with radius $R_{\\rm gal}\\simeq 10$ kpc. This gives an estimation of OB association formation rate as \\begin{equation} r_{\\rm OB}  =  \\frac{r_{\\rm II}}{N_{\\rm SN}\\,\\pi\\, R_{\\rm gal}^2}, \\end{equation} \\begin{equation} r_{\\rm OB}     \\simeq  3.2 \\times 10^{-7}{\\rm pc^{-2}\\,Myr^{-1}} \\left( \\frac{r_{\\rm II}}{0.01{\\rm yr^{-1}}}\\right) \\left( \\frac{N_{\\rm SN}}{100} \\right)^{-1} \\left( \\frac{R_{\\rm gal}}{10 {\\rm kpc}}\\right)^{-2}. \\end{equation} This indicates the two-dimensional porosity to be equal to $Q(t < 20 {\\rm Myr})\\simeq 0.5-0.6$. Thus, rather large fraction of the galactic disk, $1-\\exp{(-Q)}\\sim 40\\%-45\\%$, is covered by young ($t < 20$ Myr) superbubbles."
    },
    "9804/astro-ph9804047_arXiv.txt": {
        "abstract": "We consider magnetic field evolution of neutron stars during polar-cap accretion. The size of the polar cap increases as the field decays, and is set by the last open field line before the accretion disk. Below the polar cap we find the temperature to be so high that electron-phonon scattering dominates the conductivity. Outside the polar cap region, the temperature is such the conductivity is dominated by temperature independent impurity scattering which can be a few orders of magnitude larger than the electron-phonon conductivity. The time-scale for field decay is therefore initially given by impurity scattering dominated conductivity. When the field strength has been reduced to $\\sim 10^8 ~{\\rm gauss}$ the accretion is spherical and the time scale for field decay is given by the smaller electron-phonon scattering conductivity. The field strength is now reduced rapidly compared to before and this could be a reason for there being no pulsars known with field strengths below $10^8~{\\rm gauss}$. We also investigate the evolution of multipoles at the neutron star surface. We find that contribution from higher-order multipoles are at most 30 \\% to that of the dipole mode. ",
        "introduction": "Since the discovery of pulsars there has been much discussion of observational evidence for decay of magnetic fields in neutron stars, as well as much theoretical work. As the reviews by Lamb (1991), Chanmugam (1992), and Phinney \\& Kulkarni (1994) indicate, there is at present no consensus on the question of whether or not magnetic fields in isolated neutron stars can decay significantly.  The general view has been that the electrical conductivity of matter in the cores of neutron stars is so high that the characteristic decay time for fields generated by electrical currents in the core is greater than the age of the Universe. Recently Pethick \\& Sahrling (1995) showed that even if the conductivity in the core was small the shortest possible decay time is some two orders of magnitude longer than the decay time for configurations where the magnetic field is confined to the crust. An incorporation of general relativistic effects, Sengupta (1997), further reduces the decay rate. Still, millisecond pulsars have typical field strengths in the range $10^8~ -~ 10^{10}~{\\rm gauss}$ compared to isolated radio pulsars which have typically field strengths around $10^{12}~{\\rm gauss}$ and this indicates that during the spin-up phase the accretion process is reducing the field strength somehow.   Millisecond pulsars are generally found in binary systems where the companion star is a white dwarf with a mass less than a solar mass, $M_{\\sun}$. This system is called a Low-Mass-Binary Pulsar (LMBP) referring to the mass of the companion star. The progenitor to this system is thought to be the Low-Mass-Xray Binaries (LMXB) where a neutron star is accreting matter from a companion having a mass less than about 2 $M_{\\sun}$. The accreting matter is spinning up the neutron star to millisecond periods. For details concerning this process see reviews by Phinney \\& Kulkarni (1994) and Bhattacharya \\& van den Heuvel (1991) among others. The evolution of the binary system after the neutron star is formed either by a supernovae or tidal capture, is assumed to occur on at least two time scales when the companion star is in radiative equilibrium. At first the companion star is evolving on a nuclear time scale slowly filling its Roche lobe. When it has been filled up the matter overflows and falls onto the companion neutron star on a thermal time scale $\\tau_{th}=G M^2/R L = 5\\times10^7~ (M_{\\sun}/M)^2 ~{\\rm yrs}$, see Bhattacharya \\& van den Heuvel (1991) for details. If LMXB's are the only progenitors to LMBP's the lifetime of the LMXB, or the accretion phase of the binary system, must be of order $10^7~{\\rm yrs}$. However, by using the amount of mass needed to be accreted to spin up the neutron star to a spin period $P_i$ one finds a time scale that ranges from $10^8-10^{10}\\times (P_i/2~{\\rm ms} )^{-4/3}~{\\rm yrs}$. This discrepancy suggests that there might other progenitors to the LMBP's. For details see the review by Phinney and Kulkarni (1994). We will in this paper assume the accretion phase of progenitors to millisecond pulsars to last between roughly $10^7-10^9~{\\rm yrs}$.    Matter accreting onto a neutron star is expected to be disrupted by the magnetic field at some radius, $r_A$, sometimes called the Alfv\\'en radius. Beyond this bare description there is no generally accepted view on how or where the matter attaches to the field lines and flows to the neutron star's surface, or on the interaction of the field with the matter outside $r_A$, despite a large number of papers on the subject, see King (1995). In the case of disk accretion there are two main approaches to the problem. In one (Ghosh \\& Lamb 1978, 1979a,b; Kaburaki 1986 and Wang 1987) the field is assumed to thread a large fraction of the disk because of Kelvin-Helmholtz instabilities. The other approach (Aly 1980; Anzer \\& B\\\"orner 1980, 1983; Scharlemann 1978) assumes the disk is a perfect conductor, completely excluding the field. In both approaches the matter is often assumed to leave the disk in i a narrow transition zone at the inner edge (near $r_A$), thereafter flowing along field lines to the neutron star.  Most studies concerned with the magnetic field evolution {\\it inside} the neutron star have assumed matter to accrete onto the surface spherically, e.g. Fujimoto et al. (1984), Miralda-Escud\\'e et al. (1990), hereafter Mir90, Urpin \\& Geppert (1995), and Konar \\& Bhattacharya (1997). In this paper I consider non-spherical accretion where matter is assumed to flow onto the polar caps in a column. The cap is consequently heated up and the conductivity in the crust below the polar cap we estimate to be much smaller than the conductivity outside the accretion column. The magnetic evolution in this scenario can be divided into two stages where in stage I the global decay rate is controlled by the conductivity outside the accretion column, $\\tau_B\\sim 10^{8}-10^{10}~{\\rm yrs}$. When the field has reached a value of about $10^8 {\\rm gauss}$ the accretion is spherical and the evolution enters the second stage. Here, the whole crust is being heated up and the conductivity is dominated by electron-phonon scattering. In this stage the magnetic field decay time is a few orders of magnitude shorter than in stage I and compared to the earlier evolution the field is dissipating rapidly. I argue that this could account for the fact that no binary pulsars have been found with a magnetic field less than $10^8~{\\rm gauss}$. A similar effect where the accreted flow is pushing the field lines has already been noticed by Romani (1993).  The paper is organized as follows: section 2 contains a brief description of the model and the basic equations. In section 3 we estimate the length scale for temperature change at the accretion column-normal crust boundary to be less than a crust thickness. The time scale for the temperature to reach a stationary state is also shown to be much smaller than the magnetic field decay time scale. Therefore, we do not solve the energy equation explicitly but assume the conductivity as a function of angle to be close to a top-hat function at all times. Section 4 presents the numerical results for some initial depths of the magnetic field. ",
        "conclusions": "We have examined the influence of asymmetric accretion on the magnetic field evolution of a neutron star. The temperature structure in the crust resulting from the accretion was roughly estimated and found to vary more rapidly than the crust thickness around the edge of the accretion cap. The time scale for reaching a steady state was shown to be shorter than the time scale for magnetic field decay. Therefore, instead of doing a full blown calculation of heat conduction coupled with the magnetic field evolution we used a simple smoothed top-hat function for the temperature structure and consequently for the conductivity. The global field decay time is roughly the one given by the largest conductivity in the crust, which occurs outside the accretion cap, $\\tau_{B,i} = 5.5 \\times 10^{17}~(\\delta R_5)^2~(\\rho_{14} x/0.1)^{1/3}Z/(60~ (Q/0.01))~{\\rm s}$. However, when the field strength is down to roughly $10^8$ gauss the accretion is spherically symmetric and we have $\\tau_{B,ph} = 5.5 \\times 10^{15}~ (\\delta R_5)^2~\\rho_{14}^{7/6} ~(x/0.1)^{5/3}T_8^{-2}~ {\\rm s} \\ll \\tau_{B,i}$ and so the decay rate is much faster than initially which could provide a clue why no binary pulsars are known with field strengths less than $10^8$ gauss. Romani (1993) discussed a similar effect and found a threshold close to ours. We also found the asymmetric accretion resulted in some higher order surface magnetic multipoles and the size of these were shown to be at most 30 \\% of the surface dipole field.   {\\bf Acknowledgements :} I would like to thank Dong Lai, Lars Bildsten and Edward Brown for useful discussions. Dong Lai is especially thanked for a careful reading of the manuscript. This work was also supported in part by the U. S. National Science Foundation under grants NSF AST93-15133 and AST94-14232, by NASA under grant NAGW-1583, and by the Swedish Natural Science Research Council.  \\vfill\\eject"
    },
    "9804/astro-ph9804271_arXiv.txt": {
        "abstract": "The lack of bright host galaxies in several recently examined gamma--ray burst (GRB) error boxes suggests that the redshifts of cosmological GRBs may be significantly higher than previously hypothesized. On the other hand, the non--detection of multiple images in the BATSE 4B catalog implies an upper limit to the average redshift $\\langle z\\rangle$ of GRBs. Here, we calculate an upper limit to $\\langle z\\rangle$, independent of the physical model for GRBs, using a new statistical lensing method that removes distance ambiguities, and thus permits accurate computation of the lensing rate at high $z$. The upper limit on $\\langle z\\rangle$ depends directly on the cosmological parameters $\\Omega$ and $\\Lambda$. If there are no multiple images among the brightest 80\\% of the first 1802 bursts in the BATSE 4B catalog, then, at the 95\\% confidence level, $\\langle z\\rangle<$ 2.2, 2.8, 4.3, or 5.3 for ($\\Omega$, $\\Lambda$) values of (0.3, 0.7), (0.5, 0.5), (0.5, 0.0), or (1.0, 0.0), respectively. The 68\\% upper limit to the average redshift is comparable to or less than the median redshift of GRBs in scenarios in which the GRB rate is proportional to the rate of star formation, for any cosmology. The uncertainty in the lensing rate---arising from uncertainties in the cosmological parameters and in the number density and average velocity dispersion of galaxies---will be reduced significantly in the next few years by a new generation of experiments and surveys. Moreover, the continued increase in the number of GRBs observed by BATSE will greatly constrain their redshift distribution. ",
        "introduction": "Three decades after their discovery (\\cite{KSO73}), the physical origin of gamma--ray bursts (GRBs) remains unresolved. Recent developments, however---including the isotropy of GRBs seen by BATSE (\\cite{B96}; \\cite{THBM96}) and the detection of redshift $z$=0.835 absorption and emission lines from a possible optical counterpart to GRB 970508 (\\cite{M97}; \\cite{R98})---strongly suggest a cosmological origin for GRBs.  Initially, no--evolution fits to the log N -- log P (peak flux) distribution of BATSE GRBs suggested a typical redshift of $z\\sim 1$ for dim bursts (\\cite{F93}; \\cite{W93}), with the break in the $-3/2$~slope of the flux distribution at $P \\sim 10$ ph~cm$^{-2}$~s$^{-1}$ being interpreted as a cosmological deviation from Euclidean space at $z \\sim 1$. In addition, a number of researchers reported a factor of $\\sim 2$ ``time--stretching'' of dim bursts relative to bright ones (e.g., \\cite{Nor94}, 1995), which was thought to imply a redshift of order unity for the dim bursts. Such a redshift would imply an extremely low rate of gravitational lensing and multiple imaging of sources detected with BATSE, perhaps as low as one multiple imaging event per 200 years (\\cite{GN94}).  New lines of evidence now indicate that the typical redshift of cosmological GRBs may substantially exceed unity. Particularly suggestive evidence for high GRB redshifts comes from deep HST searches of the error boxes of several of the brightest GRBs detected by BATSE and the Interplanetary network (IPN): In five cases, no obvious host galaxies for the bursts were found, down to a limiting magnitude between 3.5 and 5.5 mag fainter than what would be expected if GRBs reside in $L_\\star$ galaxies and dim bursts are at $z\\sim 1$ (\\cite{Sch97}). This ``no--host problem'' has been analyzed further by Band \\& Hartmann (1998), and implies that if GRBs reside in normal galaxies, even the {\\em brightest} of them are at redshifts close to unity, with the faintest being at much higher redshift still. In addition, Fenimore \\& Bloom (1995) showed that the intrinsic anticorrelation between photon energy and burst duration implies that if the observed time--stretching is caused by time dilation, then the redshift of the dimmest bursts may be as large as $z\\sim 6$ (note, however, that some or all of the time--stretching may be intrinsic to the bursts; cf. \\cite{SPS97}).  These new lines of evidence have led researchers to explore scenarios in which the bursts are at much higher redshifts. One currently popular scenario, motivated in part by the assumption that cosmological GRBs involve remnants of massive stars, is that the GRB rate is proportional to the rate of (massive) star formation (\\cite{Tot97}; \\cite{WBBN98}). The star formation rate is thought to vary strongly with redshift (\\cite{Ma96}), peaking at $z \\sim 2$. If this is the case, the dimmest bursts may be at redshift $z\\sim 6$, making them the most distant objects ever detected (\\cite{WBBN98}). Other authors have noted that, if the comoving number density of GRB sources is allowed to vary with redshift, then even if the peak luminosity is fixed, the log N--log P distribution and other properties of the BATSE bursts can be fit by models with a maximum redshift up to $z\\sim$10--200 (\\cite{RHL95}).  In these high-$z$ scenarios, the expected incidence of gravitational lensing and multiple imaging of GRB sources is much higher than it is in lower-$z$ scenarios, since in general the lensing rate increases markedly with the source redshift.  The short duration of GRBs (typically tens of seconds) compared to the difference in light travel time between different ray paths (typically months, if the lens is of galactic mass; cf. \\cite{M92}) means that a multiply-imaged GRB appears as two or more separate events with identical time histories and intensities that differ only by a scale factor. The image separations, of order arcseconds, are tiny compared to BATSE location errors; thus, these GRBs would appear to come from the same location.  Two or more such events must be detected in order to identify a lensed GRB source; thus, the overall BATSE burst detection efficiency, $\\epsilon$, is of prime importance in determining the observed incidence of lensing. The average efficiency for the 4B catalog is $\\epsilon = 0.48$ (\\cite{M98}), which is 40\\% larger than the previously estimated value (\\cite{F94}). This increase implies that the expected incidence of lensing is significantly higher than was previously believed.  Lensing of GRBs was first suggested by Paczynski (1986) as a way to establish their cosmological origin. Subsequently, many authors have calculated the lensing rate of GRBs (\\cite{M92}; \\cite{BW92}; \\cite{N93}; \\cite{GN94}), assuming GRB redshifts estimated from no--evolution fits to the log N -- log P distribution, viz., $z\\lta 1$. For redshifts $z \\lta 1$, the gravitational lensing rate is reasonably well known for a given cosmology (e.g., \\cite{TOG84}; \\cite{FT91}; \\cite{F92}). At redshifts $z \\gta 1$, however, the lensing rate has been uncertain, mainly because of ambiguity in the angular diameter distance at high redshift (\\cite{F92}).  As a result, even when the cosmology and the number density and properties of lenses are fixed, the estimated lensing rate can vary by factors of several between the different prescriptions.  Motivated in part by this ambiguity, Holz \\& Wald (1998) have developed a numerical method to calculate lensing rates for a given cosmology, combining techniques from both ``ray-shooting'' and ``Swiss-cheese'' models. The approach resolves the angular-diameter distance uncertainties, and also correctly accounts for multiple lens encounters along the line of sight, allowing for an unambiguous calculation of lensing rates.  If there are no multiple images among the bursts detected with BATSE, then an upper limit to the average redshift $\\langle z\\rangle$ of GRBs can be inferred, one which depends directly on the cosmological parameters (see also Nemiroff et al.\\ 1994; Marani et al.\\ 1998; Marani 1998). Here, we set an upper limit to $\\langle z\\rangle$, independent of the physical model for GRBs, using the Holz \\&~Wald (1998) method to compute the lensing rate for a variety of cosmologies.  The plan of this paper is as follows. In \\S~2 we describe and develop the statistical lensing method, and compare its results with analytical estimates. We show that analytical estimates using the ``filled-beam\" approach are reasonably accurate for source redshifts less than $\\sim$2--3, but underestimate the true rate of lensing by tens of percent at higher redshifts. In \\S~3 we use the numerical method to compute the lensing rate as a function of redshift for several cosmologies, and compare these results with previous estimates. We also  determine upper limits on the rate of lensing and on the average redshift $\\langle z\\rangle$ of GRBs, assuming that no lensing events are present in the BATSE 4B catalog (\\cite{M98}). We discuss the implications of these results and give our conclusions in \\S~4. We follow the conventions that the Hubble constant is 100\\,$h$\\ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega$ is the present mean density in the universe in units of the closure density, and $\\Lambda$ is the present normalized cosmological constant. In a flat universe, $\\Omega+\\Lambda=1$. ",
        "conclusions": "We have exhibited a new method for calculating gravitational lensing rates, given a cosmology and a distribution and evolution of the lenses. This method, which is described more fully in Holz \\& Wald (1998), is free of the angular diameter distance ambiguity that is a feature of standard analytical methods. For the cosmologies and redshift ranges we have examined, we find that the lensing rate is approximately equal to the filled-beam lensing rate for $z<3$, but in excess of the filled-beam rate for $z>3$. We find that the lack of detected lensing in the current BATSE catalog places an upper limit to the median redshift of bursts that, at the 68\\% confidence level, is comparable to or less than the median redshift in star-formation burst models.  We now discuss these results and place them in context. In \\S~4.1 we discuss the simplifying assumptions that have been  made in our calculation. We find that these assumptions have, if anything, caused us to {\\it underestimate} the lensing rate. In \\S~4.2 we examine the uncertainties in the inputs used in our calculations, as well as the future prospects for reducing these uncertainties. In \\S~4.3 we discuss the implications of the current lack of lensing for cosmological models of gamma-ray bursts, when combined with existing lower bounds to the redshift of bursts in these models. Finally, in \\S~4.4 we provide a future outlook for the rapidly emerging importance of constraints inferred from detection or nondetection of lensing.  \\subsection{Effects of Simplifying Assumptions}  We have assumed that the angular distribution of the sources of gamma-ray bursts is uncorrelated with the angular distribution of the lenses. This is well justified, because for sources at low redshift the lenses that contribute most to the lensing rate are at about half the redshift to the source, and for sources at high redshift the dominant lenses are at a redshift of approximately unity (see \\S~2; see also, e.g., Turner et al.\\ 1984; \\cite{M92}). Hence, in a cosmological model the sources of gamma-ray bursts are separated from lenses by hundreds of megaparsecs to gigaparsecs, and therefore it is extremely unlikely that the sources and lenses are angularly correlated. We have also assumed that the phase of the orbit of BATSE is uncorrelated between the separate images of the burst, and hence that the joint probability of two images being observable is just the square of the BATSE efficiency, $0.48^2 \\approx 0.23$. This assumption is reasonable if the time delay between images is significantly larger than the BATSE orbit time of $\\sim 5000$ seconds. The time delay for a mass $M$ is of order $10^{-5}(M/M_\\odot)$ seconds (see, e.g., Blandford \\& Narayan 1992), and hence the assumption of uncorrelated phases is good for masses $M\\gta 10^9\\,M_\\odot$. This includes almost all the effective lensing mass of galaxies, and thus this assumption is also unlikely to affect the calculated lensing rates significantly.  We may have underestimated the lensing rate by assuming that only galaxies---which we have modeled as singular isothermal spheres---contribute to the lensing rate. Nemiroff et al.\\ (1993) have noted that point masses, for example supermassive black holes, could also in principle contribute. To produce gamma-ray burst lensing of the type that we consider here, in which the lensing event appears as two separate bursts, the mass of such a lensing black hole must be $M\\gta 10^8\\,M_\\odot$, because smaller masses would produce time delays less than $\\sim$1000 seconds, so that overwriting or readout time would tend to prevent BATSE from detecting two separate bursts (the effect of lensing by lower-mass black holes may, however, be detectable using techniques such as autocorrelation analysis; see Nemiroff et al.\\ 1994, 1998). If, however, the total lensing rate were dominated by very massive black holes, $M\\gta 10^{10}\\,M_\\odot$, then these objects would also produce detectable image separation of lensed quasars or galaxies (the angular separation is of order $3(M/M_\\odot)^{1/2} (D/1 {\\rm Gpc})^{-1/2}\\,\\mu$arcsec for an Einstein ring, which is 0.3\" for $M=10^{10}\\,M_\\odot$ and $D$=1 Gpc). Hence, the amount by which we have underestimated the lensing rate depends on the number of point masses in the relatively narrow range $10^8$--$10^{10}\\,M_\\odot$, which is the only mass range that could avoid detection in quasar lensing surveys and yet produce separately detected bursts.  We may also have underestimated the lensing rate by only keeping track of the total number of images, not whether there are, e.g., two or four images in a particular event. As pointed out by Grossman \\& Nowak (1994), this tends to underestimate the {\\it detectable} rate of lensing because when there are several images it is more probable that at least two of them are observed by BATSE: assuming a 48\\% efficiency for any particular image, the probability of observing at least one pair rises from 23\\% when there are two images to 66\\% when there are four images. Grossman \\& Nowak (1994) estimate that, for galaxy ellipticities typical of galaxy surveys, the maximum possible enhancement of the lensing detection rate is $\\sim$30\\%, but that the overall enhancement is likely to be smaller.  A third effect that may enhance the rate of lensing is magnification bias, which is an effect first discussed extensively by Fukugita \\& Turner (1991) in the context of quasar lensing. Magnification bias occurs because sources that would have been undetectable are made visible by lensing, and hence a magnitude-limited sample contains an enhanced incidence of lensing. This can be especially important if the number of sources at a given flux rises steeply at the faint end, as it does in quasars. To detect multiple images and thus confirm strong lensing, it is necessary to detect the fainter image as well as the stronger image. For persistent sources such as quasars, this can be done by following up broad, shallow surveys by deep pointings, so that essentially all of the secondary images are detectable (e.g., as in the Hubble Snapshot Survey [Bahcall et al. 1992; Maoz \\etal\\ 1992, 1993a,b]). For transient sources such as gamma-ray bursts, deep follow-ups are not always possible, and hence for the secondary image to be detectable it must have a brightness in excess of the threshold of the original survey (i.e., in the case of gamma-ray bursts, the BATSE sample). This effect, plus the flatness of the faint end of the gamma-ray burst log N -- log P curve, suggests that the magnification bias for gamma-ray bursts is probably less than the magnification bias for quasars (see also \\cite{GN94} for a discussion of this point). Nonetheless, magnification bias for gamma-ray bursts could in principle increase significantly the expected lensing rate, and hence decrease the upper limits on $\\langle z\\rangle$, compared to the conservative estimates here.  We may underestimate the lensing rate because we have neglected the clustering of galaxies, the last of our assumptions. The enhanced gravitational potential in clusters tends to increase the convergence of null geodesics and hence increase the cross section for strong lensing (see also Holz \\& Wald 1998). Numerical estimates suggest that the lensing rate for galaxies in a cluster will be increased by $\\sim$10--20\\% by this effect, but because only $\\sim$10\\% of galaxies are in clusters this effect is unlikely to increase the overall lensing rate significantly.  Finally, however, we may have overestimated the rate of lensing by ignoring evolution in the properties of the lensing galaxies. If the comoving density in lensing galaxies was less in the past than it is now, or if these galaxies were less massive than they are today, high-redshift lenses would contribute less to the lensing rate than we have assumed. Note, however, that the peak in the lens redshift distribution is at about unity (Fig.~2), when galaxies had properties very similar to their current properties. Note also that observations of lensing of quasars (see below), which have a typical redshift similar to the proposed typical redshift $z\\sim 2$ of GRBs, suggest that if there are evolutionary effects then these effects have not had a dominant effect on the lensing rate.  To summarize, the net effect of our simplifying assumptions is likely to be that we underestimated the lensing rate by $\\lta$10\\%, and hence our results our conservative in that they give a slightly high upper bound to $\\langle z\\rangle$.  \\subsection{Uncertainties in Input Parameters}  In addition to the simplifying assumptions discussed above, our calculations are clearly dependent on the input values we have assumed. One input is the dimensionless parameter $F$, which is proportional to the product of the number density of galaxies and the fourth power of their average velocity dispersion. The lensing rate scales linearly with $F$.  We have used $F$=0.1, which is consistent with the recent estimates from the Century Survey (Geller et al.\\ 1997). Other estimates range from $F$=0.05 to $F$=0.15, and if these other estimates are used then the range of possible maximum redshifts is increased for a given cosmology. Fundamental cosmological parameters, another input to our calculations, are also uncertain. A standard Einstein-De Sitter universe ($\\Omega$=1, $\\Lambda$=0) gives the lowest lensing rate, and a $\\Lambda$-dominated universe gives the highest.  Current observation of quasar lensing puts constraints on the allowed combinations of $F$ and cosmology. For example, five of the 351 quasars in the HST Snapshot Survey (Maoz et al.\\ 1993b) were lensed. A comparison of this rate with the rate expected for models is complicated by effects such as magnification bias (see above). However, the high average redshift of the quasars ($z\\sim 2$) means that many of the issues affecting lensing of gamma-ray bursts, such as evolution of clustering, also affect the lensing of quasars. Hence, quasar lensing statistics have bearing on estimated GRB lensing rates, and will become even more relevant as more extensive surveys are done.  In addition, all of these uncertainties will be diminished greatly by the data that emerge from the plethora of satellites and surveys planned for the next few years. The number density and velocity distribution of galaxies (and hence $F$) will be determined with unprecedented accuracy by surveys such as 2dF (Colless 1998) and the Sloan Digital Sky Survey (Margon 1998), the latter of which is expected to see first light in 1998. Data from these surveys will also provide valuable information about  $\\Omega$, $\\Lambda$, clustering, and the lensing rate of quasars. Complementary information about cosmological parameters will be extracted from high-redshift supernova surveys (Perlmutter \\etal\\ 1998) and from data gathered by the many upcoming microwave background experiments, such as MAP (Wang, Spergel, \\& Strauss 1998), TOPHAT (Martin et al.\\ 1996), and PLANCK (see, e.g., Bond, Efstathiou, \\& Tegmark 1997 for a discussion of the constraints). The evolution of galaxy clustering is already being inferred from cluster surveys, and AXAF observations of X-ray emission from clusters of galaxies is expected to greatly enhance this understanding. Hence, one side effect of the coming data-rich era in cosmology is that calculations of lensing rates will be much less uncertain, and therefore upper limits to the redshift of a population such as gamma-ray bursts will be far more secure.  \\subsection{Current Implications for Cosmological GRB Models}  The redshift limits presented in this paper have important consequences for cosmological models of GRBs. Lower limits on the redshift of bright GRBs are becoming stronger as the number of detected bursts rises, both because of the no-host problem (Schaefer et al.\\ 1997) and because GRBs do not show any evidence of large-scale clustering (Lamb \\& Quashnock 1993; Quashnock 1996). These limits suggest that if the sources of GRBs are in or near galaxies, the median redshift of the bursts detected with BATSE may be significantly greater than unity. However, the upper limits to the redshift  of GRBs that follow from the lack of lensing are beginning to make such  high-redshift populations less appealing. For example, the lack of lensing in the current BATSE catalog strongly rules out the \\hbox{$z\\sim 10$--200} models that were heretofore consistent with the properties of the BATSE bursts (Rutledge et al.\\ 1995). Also, models in which the burst rate is proportional to the star formation rate (Totani 1997; Wijers et al.\\ 1998), and for which $\\langle z\\rangle\\sim 2.2$, are beginning to be constrained strongly by the lack of lensing: Our 68\\% upper limit on $\\langle z\\rangle$ (eq.~[6]) is less than 1.9, and probably close to 1.0. If the redshift upper limits are reduced further, it may be necessary to postulate either a very short interval in which GRBs were produced, with an intrinsic brightness distribution that matches the observed log N -- log S distribution, or a population that is unassociated with any known population of objects, so that the lower limits to redshift do not apply.   \\subsection{Future Outlook and Summary}  In the next few years gravitational lensing is likely to play a role  of rapidly increasing importance in constraining cosmological models of gamma-ray bursts. As discussed in \\S~4.2, the uncertainties in the calculation of the expected lensing rate will be diminished greatly, and hence the limits will be robust. In addition, as BATSE continues to detect bursts the statistics will steadily improve, and in the current high-redshift models lensing is to be expected in the next few years. In five years, the sample will be roughly double the $\\sim$1800 bursts considered here, and therefore if lensing continues to be absent, upper limits on the redshift will become extremely restrictive. At the same time, the improved statistics will continue to increase the lower limits on redshift, if there is no apparent large-scale clustering (Lamb \\& Quashnock 1993; Quashnock 1996). Moreover, the accurate positional estimates of BeppoSAX,  HETE II, and the Fourth Interplanetary Network (IPN) are likely to provide $\\sim$100 error boxes of area a few square arcminutes or less in the next five years.  If instead lensing {\\it is} detected, there will be many important  consequences. For one, the cosmological origin of a particular GRB will be established from gamma-ray data for the first time (as opposed to from the afterglow, as may be the case for GRB~970508 [\\cite{M97}]). The approximate mass of the lens, and hence its luminosity, can be estimated from the time delay between images. The sharp peak in the probable redshift of lenses (see \\S~2) would then give a reasonably accurate estimate of the flux of the lens. If the GRB can be reasonably well-localized (e.g., with the IPN), this will allow a search for the lens. If the lens is found, the GRB source must essentially be angularly coincident with it, because the deflection angle is only about 1\", and thus counterpart searches will be facilitated greatly. Knowledge of the redshift of the lens, combined with $H_0$ and an independent estimate of the mass of the lens (e.g., from its luminosity or velocity dispersion) will allow use of the time delay to estimate the redshift of the source. The quantifiable diversity of gamma-ray burst light curves (\\cite{MN98}; \\cite{Mar98}) means that false positives are unlikely, and hence even a single lensing event will generate a wealth of data.  In conclusion, the increased estimates of the redshift of gamma-ray bursts that were forced by the no-host problem, in addition to the greatly improved statistics of bursts and the increase in the estimate of the BATSE burst detection efficiency from 34\\% to 48\\%, mean that gravitational lensing limits are now playing a prominent role in constraining gamma-ray burst models. The role of lensing will become dramatically more important and robust in the next five years."
    },
    "9804/astro-ph9804101_arXiv.txt": {
        "abstract": "The blazar Mkn 421 has been observed, as part of the  AO1 Core Program, five times from 2 to 7 May 1997. In the LECS+MECS energy band the spectrum shows convex curvature, well represented by a broken power--law. Flux variability (more than a factor 2) has been detected over the entire 0.1--10 keV range, accompanying which the spectrum steepens with the decrease in intensity. Mkn 421 has also been detected with the PDS instrument. Our preliminary analysis indicates that the PDS spectrum lies significantly above the extrapolation from the MECS, suggesting a contribution from a flatter high energy component. ",
        "introduction": "Mkn 421 is one of the best known and studied BL Lac objects. It shows optical polarization, flat radio spectrum, and large variability, characteristics of the blazar class. It is bright in X--rays, where it shows prominent flares accompanied by significant spectral changes \\cite{takahashi}.  Among GeV--emitting blazars, Mkn 421 is unique in being the first object in which the $\\gamma$--ray emission is detected to extend up to TeV energies (\\cite{punch,krennrich}), at a level allowing detailed spectral and variability studies in the broadest available range of frequencies. Its GeV (EGRET) $\\gamma$--ray emission connects smoothly with the E$>$0.5 TeV spectrum (e.g.~\\cite{macomb}). On the contrary the X--ray spectrum, up to 10 keV, did not show any hint of the onset of the inverse Compton component responsible for GeV and TeV emission. It was then of great interest to take advantage of the full capabilities of {\\it Beppo}SAX, enabling a spectral coverage up to $\\gta$ 200 keV, to look for it. At the same time the fact that in the X--ray band we are observing the emission from the highest energy tail of the emitting--particle distribution could provide important clues on particle acceleration and cooling mechanisms.  The Mkn~421 observations discussed here are part of the AO1 Core Program dedicated to bright blazars. The data reduction presented here has been generally performed with software released {\\it before} September 1997. Data reduction with the updated software for all the on board instruments, and a more detailed analysis, will be presented in forthcoming paper with all appropriate references.  \\begin{figure}[t] \\vspace{9pt} \\psfig{file=lecs_2mecs_lc.ps,width=7.5truecm,rheight=5.3truecm} \\caption{\\small\\sf Rebinned light curve (4000~s bins) of {\\it Beppo}SAX data, LECS, and MECS divided in two energy bands [1.5 -- 4], [4 -- 10] keV.} \\label{fig:light_curve} \\end{figure}   \\begin{table*}[t] \\setlength{\\tabcolsep}{0.7pc} \\newlength{\\digitwidth} \\settowidth{\\digitwidth}{\\rm 0} \\catcode`?=\\active \\def?{\\kern\\digitwidth} \\caption{{\\it Beppo}SAX Observations Log} \\label{tab:log} \\begin{tabular*}{\\textwidth}{@{}l@{\\extracolsep{\\fill}}ccccc} \\hline {Pointing} & {LECS } & {LECS } & {MECS } & {MECS } & {MECS } \\\\ {Start Date} & {exp. time} & {[0.1--4 keV]} & {exp. time} & {[1.5--10 keV]} & {[$<$4/$>$4 keV]} \\\\ { } & {(ksec)} & {(cts/s)} & {(ksec)} & {(cts/s)} & {(cts/s)} \\\\ \\hline  { 2/V/1997 @ 04:10 } & {  4.4 } & {$ 2.806 \\pm 0.020 $} & { 11.4 } & {$ 2.862 \\pm 0.016 $} & { (2.14/0.72) } \\\\ { 3/V/1997 @ 03:24 } & {  4.3 } & {$ 1.748 \\pm 0.020 $} & { 11.7 } & {$ 1.694 \\pm 0.012 $} & { (1.28/0.41) } \\\\ { 4/V/1997 @ 03:25 } & {  4.9 } & {$ 1.362 \\pm 0.017 $} & { 12.2 } & {$ 1.027 \\pm 0.009 $} & { (0.81/0.22) } \\\\ { 5/V/1997 @ 03:32 } & {  4.9 } & {$ 1.823 \\pm 0.029 $} & { 11.9 } & {$ 1.522 \\pm 0.012 $} & { (1.52/0.35) } \\\\ { 7/V/1997 @ 04:47 } & {  6.0 } & {$ 1.612 \\pm 0.016 $} & { \\nodata } & { \\nodata } & { \\nodata } \\\\ \\hline \\end{tabular*} \\end{table*} ",
        "conclusions": "{\\it Beppo}SAX dat aof Mkn 421 shows interesting variability both in flux and in spectral shape, with a marked softening corresponding to decreasing brightness. This kind of spectral variability behavior is well known in the X--ray band for sources of the class of Mkn~421, the so called High-Frequency-Peaked BL Lacs (HBL). In general in blazars at energies just above the synchrotron peak it the relationship {\\it harder-when-brighter} holds and is generally interpreted in terms of injection of fresh electrons in the highest energy end of a single population.  On May 2$^{\\rm nd}$ the peak of the synchrotron component could possibly fall in/or just below the LECS band, while in the lower state of May 4$^{\\rm th}$ this is not longer true.  Moreover, preliminary analysis of PDS data suggests the presence of a deviation from the continuously downward curvature for E $> 10$ keV, possibly being the signature of the onset of a different harder spectral component.   Further temporal and spectral analysis, also with comparison of multifrequency data and comparison with model predictions is in progress."
    },
    "9804/astro-ph9804337_arXiv.txt": {
        "abstract": "A statistical study of global galaxy parameters can help to improve our understanding of galaxy formation processes.  In this paper we present the analysis of global galaxy parameters based on optical and near-infrared observations of a large sample of edge-on disc galaxies. \\\\ We found a correlation between the ratio of the radial to vertical scale parameter and galaxy type: galaxies become systematically thinner when going from S0's to Sc's, whereas the distribution seems to level off for later types.  \\\\ The observed scale length ratios (and thus the radial colour gradients) largely represent the galaxies' dust content.  On average the colour gradients indicated by the scale length ratios increase from type Sa to at least type Sc.  For galaxy types later than Sc, the average colour gradient seems to decrease again.  \\\\ The distribution of {\\it K}-band (edge-on) disc central surface brightnesses is rather flat, although with a large scatter.  However, the latest-type sample galaxies ($T > 6$) show an indication that their average disc central surface brightnesses may be fainter than those of the earlier types.  This effect is probably not the result of dust extinction. ",
        "introduction": "A study of the statistical properties of highly inclined, or ``edge-on'' galaxies benefits greatly from the special orientation with respect to the line of sight of such galaxies.  Observations of edge-on galaxies provide us with direct measurements of the luminosity and colour distributions both perpendicular to the galaxy planes and along the galaxies' major axes at various heights above the plane.  Indirectly, these luminosity distributions can be related to the galaxies' density distributions and thus their global structure.  Moreover, in-depth knowledge of the dust distribution, and hence the optical depth of galaxies, is important for our understanding of galaxy evolution.  \\subsection{The flattening of exponential discs} \\label{ratio.sec}  A major advantage of studying highly inclined galaxies is that one can determine their radial and vertical scale parameters directly and independently, since the dependence of these parameters on inclination is smallest for the highest inclinations (e.g., van der Kruit \\& Searle 1981a; de Grijs et al.  1997).  These scale parameters provide us with information about the intrinsic shape of galaxy discs, i.e., their flattening, in a more direct way than the canonical axis ratios. Moreover, since the vertical scale height, $z_0 = 2 h_z$ (where $h_z$ is the exponential scale height), is to first order independent of radius (e.g., van der Kruit \\& Searle 1981a,b, 1982a; Kylafis \\& Bahcall 1987; Shaw \\& Gilmore 1990; Barnaby \\& Thronson, Jr.  1992; but see de Grijs \\& Peletier 1997), the radial to vertical scale parameter ratio, $h_R/z_0$, can often be determined more accurately than the major to minor axis ratio.  By studying the scale parameter ratio statistically, we may be able to put constraints on the disc formation processes as well as on the stability of galaxy discs (e.g., Bottema 1993).  When considering the physical processes that determine the scale parameters one does not immediately expect a strong correlation between scale length and scale height.  The scale height is likely determined by the internal, secular evolution of the stellar velocity dispersion (e.g., van der Kruit \\& Searle 1981a; Carlberg 1987), whereas the scale length is basically the result of the composition of the protogalaxy (Fall 1983; van der Kruit 1987).  However, one might expect that in a larger disc, with a greater rotation velocity, the heating of the disc stars may be more violent, thus resulting in a larger scale height.  Therefore, one may expect a correlation between the rotation velocity (which can be related directly to the scale length) of a galaxy disc and the scale height, although the precise dependence is yet unknown (see, e.g., Bottema 1993).  Thus, statistics on the ratio of scale length to scale height can be expected to give information on the importance of the formation processes in disc galaxies with different properties.  Moreover, once the $h_R/z_0$ ratio is known, one may be able to determine the (theoretical) maximum rotation of a disc from measurements of the vertical disc dispersion (Bottema 1993).  Therefore, a statistical treatment of the scale parameter ratio may put general constraints on both the kinematical properties and the global stability of galaxy discs.  Bottema (1993) predicts that a constant value for the $h_R/z_0$ ratio leads to a more or less constant mass-to-light ratio of the old disc, $(M/L)_B$, under the assumption that we are dealing with exponential, locally isothermal discs with a constant ratio of vertical to radial velocity dispersion.  On the other hand, if we assume a linear relationship between the old-disc absolute luminosity and the vertical velocity dispersion, Bottema (1993) shows that, for a constant $(M/L)_B$, the $h_R/z_0$ ratio decreases rapidly from faint galaxies to a constant level for normal and bright galaxies.  Thus, in general, the observed velocity dispersions imply that a constant old-disc mass-to-light ratio results in an approximately constant scale parameter ratio, whereas a constant scale parameter ratio also leads to a mass-to-light ratio that is, to first order, constant.  In fact, these predictions imply that all galaxy discs are governed by equal mass-to-light ratios in the old stellar populations, assuming that all galaxy discs have approximately the same colour (Bottema 1993).  However, the assumption of a constant and equal mass-to-light ratio of the old-disc population in disc galaxies is probably not physically realistic, considering the range of colours observed within and among galaxies (e.g., de Jong 1996c, and references therein).  Therefore, the predicted relationships should be treated with caution and only be used as general guidelines.  \\subsection{Colour gradients as diagnostics}  Broad-band colours are relatively easy to obtain and are therefore the most widely used colour diagnostics to date.  They immediately reveal the approximate nature of a galaxy, which is to first order determined by its dominant stellar population and dust content.  Although for the detailed analysis of galaxy luminosity and colour profiles one needs to adopt {\\it a priori} assumptions concerning the evolutionary stellar population synthesis, the initial mass function, the metallicity and the star formation history, as well as about the dust geometry and its characteristics, de Jong (1996c) shows that the colours formed from different broad-band combinations correlate strongly, which indicates that these colours are probably caused by the same physical process.  Therefore, broad-band colours can be used as indicators of changes in the gross properties of galaxies (e.g., changes in metallicity and/or dust contamination).  All systematic colour differences induced by stellar population changes and metallicity gradients are generally considerably smaller than the reddening due to dust, however.  \\subsubsection{Radial colour gradients in edge-on disc galaxies} \\label{edgeongrad.sect}  In contrast to the large number of studies of radial colour gradients in moderately inclined and face-on spiral galaxies (e.g., de Jong 1996c, and references therein), the colour behaviour of highly inclined and edge-on galaxies has not received much attention.  In highly inclined galaxies, the study and interpretation of intrinsic colour gradients is severely hampered by the presence of dust in the galaxy planes, which causes the dust lane to appear as a red feature in vertical colour profiles (e.g., Hamabe et al.  1979; Hegyi \\& Gerber 1979; van der Kruit \\& Searle 1981b; Jensen \\& Thuan 1982; de Grijs et al.  1997).  In individual edge-on galaxies, it is generally found that the colours along the major axes, i.e., the locations of the dust lanes, remain nearly constant (e.g., Sasaki 1987; Wainscoat et al. 1990; Aoki et al. 1991; Peletier \\& Balcells 1997), although in most cases the outermost disc regions tend to be slightly bluer on the major axis (e.g., Sasaki 1987), which may be explained in terms of an increasingly metal-poor population or a decreased amount of dust at larger galactocentric distances.  Generally, as the height above the dust lane and its embedded young disc increases, the radial colour gradients become small or statistically insignificant (e.g., Hamabe et al. 1979, 1980; van der Kruit \\& Searle 1982a,b; Jensen \\& Thuan 1982; Peletier \\& Balcells 1997).  \\subsubsection{Colour gradients from scale length ratios} \\label{colgrads.sect}  Since the dust influence varies as a function of passband, scale length ratios could be used as a diagnostic to estimate colour gradients and the dust content of a given galaxy.  Evans (1994) studied the effects of dust on the stellar scale length as a function of wavelength, under the assumption that the resulting scale length differences are solely due to dust absorption.  His models predict that these differences are small, at least for face-on galaxies, on the order of the observational uncertainties, and even smaller for galaxies with a prominent bulge component.  According to his models, if the scale height ratio between dust and stars is $\\sim 0.5$ (Peletier \\& Willner 1992; Evans 1994), Evans' (1994) models exclude face-on galaxies with $h_B/h_H \\approx 2$.  On the other hand, larger ratios can be obtained if a galaxy is inclined with respect to the line of sight.  The measurement of blue to red scale length ratios alone will not unambiguously reveal the dust content of a given galaxy, because any deviation from unity can equally well be explained by an intrinsic colour gradient, especially for face-on galaxies (Byun et al. 1994). ",
        "conclusions": ""
    },
    "9804/astro-ph9804011_arXiv.txt": {
        "abstract": "In this paper we demonstrate for the first time the connection between the spatial and temporal progression of star formation and the changing locations of the very dense regions in the gas of a massive disk galaxy (NGC~1144) in the aftermath of its collision with a massive elliptical (NGC~1143). These two galaxies form the combined object Arp 118, a collisional ring galaxy system. The results of 3D, time-dependent, numerical simulations of the behavior of the gas, stars, and dark matter of a disk galaxy and the stars and dark matter in an elliptical during a collision are compared with multiwavelength observations of Arp 118. The collision that took place approximately 22 Myr ago generated a strong, non-linear density wave in the stars and gas in the disk of NGC~1144, causing the gas to became clumped on a large scale. This wave produced a series of superstarclusters along arcs and rings that emanate from the central point of impact in the disk. The locations of these star forming regions match those of the regions of increased gas density predicted the time sequence of models.  The models also predict the large velocity gradients observed across the disk of NGC 1144. These are due to the rapid radial outflow of gas coupled to large azimuthal velocities in the expanding ring, caused by the impact of the massive intruder. ",
        "introduction": "Optically-identified, collision-produced 'ring galaxies' usually display vigorous star formation in an expanding ring or arc of high density gas (e.g., Joy \\& Harvey, 1987; Joy \\& Ghigo, 1988; Appleton \\& Marston, 1997; see Appleton \\& Struck-Marcell, 1996). This recent star formation often dominates the appearance of these galaxies, with H$\\alpha$ emission tracing the giant HII regions (e.g., Hippelein, 1989a) and outshining the nucleus unless the latter is active. Near-infrared images of these systems show that the older stellar disk population has also been swept into an outwardly expanding, strong density wave.  The density disturbance produced by a collision through (and roughly perpendicular to) a galaxy's disk propagates outward through a rotating disk of gas and stars, which is itself at first contracting and then expanding as the collision proceeds. This superposition of material motions and wave propagation produces a pattern of both closed loops and open-ended arcs of relatively high density gas. Shocks can occur in these regions because of the high relative velocities that are produced in the flows. The higher density features are well delineated in observations of these systems and in the combined N-body/hydrodynamic models that we and others have produced. The models can be exploited to further our understanding of the star formation that is triggered by a galaxy collision, a process that has likely occurred in many galaxies over their lifetimes and may have been very frequent earlier in the universe (see Lavery {\\it et al.}, 1996).  Here we compare our numerical models of these types of galaxy systems with observational data on Arp 118, one particular IR-luminous, gas-rich example of a collisionally produced ring galaxy. This system consists of a strongly distorted disk galaxy and an elliptical in close proximity. Hippelein (1989a) had difficulty in explaining the extreme velocity gradient and complex morphology in the Arp 118 system using a his simple picture of a collision between a gas-rich spiral and an elliptical. However, Gao {\\it et al.} (1997) have observed that Arp 118 contains a large amount of molecular gas distributed exclusively along the ring formations (the first high resolution CO observations ever made of a ring galaxy) , and that the velocity structure in this gas is kinematically consistent with the simple collisional model.  In this paper, we show that, using fully dynamical, 3D models, we can reproduce the morphology of the disturbed disk galaxy in the pair and the approximate relative positions of the two galaxies. The 'best-fit' model for the Arp 118 system was chosen from a grid of simulations produced by Gerber, Lamb, \\& Balsara (1996). These simulations explore the results of face-on collisions (collisions parallel to the spin-axis of the disk) between an elliptical galaxy and a disk galaxy. This dynamical treatment confirms the correspondence between the models and the observed velocity structure. We use the chosen simulation to constrain the timescales for star formation in Arp 118 and to explore the history of the collision by comparing the predicted results of such a collision with current observations. ",
        "conclusions": "The simulation presented here is comprised of a sequence of models that follows the evolution of a pair of galaxies through and after a collision. One of the later models closely matches the observed morphology of the CO component in Arp 118 and its velocity field. Earlier models in the sequence provide a good match with the present radio continuum and H$\\alpha$ emission, indicating that the sequence of post-collision star formation can be traced and timed by comparison to simulations of encounters between two massive galaxies with masses similar to those observed.  Stars, particularly massive stars, are formed in the cores of giant molecular clouds in the highest density regions. Both multiwavelength observations and our numerical models of Arp 118 indicate that strong shocks in the gas together with large increases in the gas volume density are associated with star formation over volumes of 1 kpc$^3$. The observed morphology of the regions of dense gas and the clustering observed in the stars emitting H$\\alpha$, which were formed in the gas since the collision, suggest that the observed clumping of the young stars results from a clumping of the densest gas on the same scale. The recent work by Marston \\& Appleton (1995) and Appleton \\& Marston (1997) also provides evidence that the clumping observed in the optical images of collisionally produced ring galaxies is not due to patchy dust obscuration, because the same clustering is also observed in the near infrared. Gas clumping on this same scale is found in the numerical simulations, suggesting that there is a global explanation for the observed morphology of the dense gas and the resulting giant stellar formations in these systems. The simulations show a relatively small perturbation (clumping) in the density of stars. We therefore predict that the distribution of the old stars in these systems will be observed to be smoother than that of the gas, although the models show that the stars are driven into a wide ring, and sometimes even a second, inner ring, by the collision (see Lamb {\\it et al.} 1998).  This study of the Arp 118 system demonstrates that a careful comparison between high resolution observations and detailed models can yield insight into the sequence of star formation that takes place in a gas-rich galaxy after a major collision. The intensity and location of the starburst at any particular epoch will depend upon the speed with which density waves are propagating through the expanding disk. Such quantities can now be predicted quite accurately from current models of colliding galaxies. Thus global star formation (on the scale of several hundred parsecs) as it occurs in these systems at the current epoch can be investigated more thoroughly than previously. The rate of galaxy collisions in the past was larger than it is today, due to the greater overall density, thus a considerable portion of the star formation that took place in young disk galaxies at earlier epochs was likely triggered by galaxy collisions. We expect, therefore, that studies like the one reported here will help in understanding this earlier star formation and its current consequences."
    },
    "9804/astro-ph9804227_arXiv.txt": {
        "abstract": "A comparison is made between the properties of CAL 83, CAL 87, RX J0513.9$-$6951, 1E 0035.4$-$7230 (SMC 13), RX J0019.8+2156, and RX J0925.7$-$4758, all supersoft X-ray binaries.  Spectra with the same resolution and wavelength coverage of these systems are compared and contrasted.  Some new photometry is also presented.  The equivalent widths of the principal emission lines of H and He II differ by more than an order of magnitude among these sources, although those of the highest ionization lines (e.g. O VI) are very similar.  In individual systems, the velocity curves derived from various ions often differ in phasing and amplitude, but those whose phasing is consistent with the light curves (implying the lines are formed near the compact star) give masses of $\\sim$1.2M$_{\\odot}$ and $\\sim$0.5M$_{\\odot}$ for the degenerate and mass-losing stars, respectively.  This finding is in conflict with currently prevailing theoretical models for supersoft binaries.  The three highest luminosity sources show evidence of ``jet\" outflows, with velocities of $\\sim$1--4$\\times$10$^3$ km s$^{-1}$.  In CAL 83 the shape of the He II 4686\\AA\\ profile continues to show evidence that these jets may precess with a period of $\\sim$69 days. ",
        "introduction": "The close-binary ``supersoft sources\" (SSS) are now recognized as a distinct class of very luminous (L$_{bol}\\geq10^{38}$ erg s$^{-1}$) X-ray sources characterized by extremely soft X-ray spectra with little or no radiation above $\\sim$0.5 keV (e.g. Tr\\\"umper et al.\\ 1991).  Several reviews of the observational properties of these sources have recently been published (e.g. Hasinger 1996, Greiner 1996, Kahabka \\& van den Heuvel 1997).  All SSS appear to have high mass-accretion rates and exhibit long-term X-ray and optical variability which are thought to reflect variations in the rate of mass transfer.  In addition, some SSS show evidence of collimated outflows or ``jets\" (Crampton et al.\\ 1996, hereafter CHC96; Southwell, Livio \\& Pringle 1997).  Van den Heuvel et al.\\ (1992) suggested that the X-ray properties of SSS are best explained by a model involving steady nuclear burning on the surface of a white dwarf accreting material at the Eddington rate.  Many observations appear to support this model (Greiner 1996), although alternative interpretations (e.g. Kylafis \\& Xilouris 1993) have not yet been ruled out.  During a 1996 November CTIO observing run we obtained spectra and some photometry of six close-binary supersoft sources.  One of these lies in the Small Magellanic Cloud, 1E 0035.4$-$7230 (hereafter SMC 13).  CAL 83, CAL 87, and RX J0513.9$-$6951 (hereafter RX J0513) are all members of the Large Magellanic Cloud, while RX J0019.8+2156 (hereafter RX J0019) and RX J0925.7$-$4758 (hereafter RX J0925) are galactic systems.  Since these objects were observed with the same spectrographic configuration, intercomparison of their spectra is very straightforward.  In addition, we present previously unpublished photometry for CAL 83 and a few observations of RX J0513 and RX J0925.  New data for CAL 87 is being published in a separate paper (Hutchings et al.\\ 1998).  Long-term spectroscopic and photometric monitoring of these sources is important since the SSS exhibit significant variations over timescales of months and years.  A summary of the properties of these six supersoft binaries is given in Table 1, where they are listed in order of decreasing orbital period.  The bolometric luminosity (L$_{bol}$) listed is the average value given in Greiner's catalog (1996); it depends strongly on the assumed model, adopted distance, and amount of absorption assumed. ",
        "conclusions": ""
    },
    "9804/astro-ph9804157_arXiv.txt": {
        "abstract": "The scattering diameters of \\sgra\\ and several nearby OH masers ($\\approx 1\\arcsec$ at 1~GHz) indicate that a region of enhanced scattering is along the line of sight to the Galactic center.  We combine radio-wave scattering data and free-free emission and absorption measurements in a likelihood analysis that constrains the following parameters of the GC scattering region: The GC-scattering region separation, $\\delgc$; the angular extent of the region, $\\psi_\\ell$ and $\\psi_b$; the outer scale on which density fluctuations occur, $l_0$; and the gas temperature, $\\te$.  The maximum likelihood estimates of these parameters are $\\delgc = 133_{-80}^{+200}$~pc, $0.5\\arcdeg \\le \\psi_\\ell \\lesssim 1\\arcdeg$, and $(l_0/1\\,\\mathrm{pc})^{2/3}\\te^{-1/2} = 10^{-7 \\pm 0.8}$.  The parameter $\\psi_b$ was not well constrained and we adopt $\\psi_b = 0\\fdg5$.  The close correspondence between $\\delgc$ and $\\psi_\\ell\\dgc$ suggests that the scattering region encloses the \\hbox{GC}.  As host media for the scattering, we consider the photoionized surface layers of molecular clouds and the interfaces between molecular clouds and the $10^7$~K ambient gas.  We are unable to make an unambiguous determination, but we favor the interface model in which the scattering medium is hot ($\\te \\sim 10^6$~K) and dense ($n_{\\mathrm{e}} \\sim 10$~cm${}^{-3}$).  The GC scattering region produces a 1~GHz scattering diameter for an extragalactic source of 90\\arcsec, if the region is a single screen, or 180\\arcsec, if the region wraps around the GC, as appears probable.  We modify the Taylor-Cordes model for the Galactic distribution of free electrons in order to include an explicit GC component.  We predict that pulsars seen through this region will have a dispersion measure of approximately $2000$~pc~cm${}^{-3}$, of which approximately 1500~pc~cm${}^{-3}$ arises from the GC component itself.  We stress the uniqueness of the GC scattering region, probably resulting from the high-pressure environment in the \\hbox{GC}. ",
        "introduction": "\\label{sec:gc.intro}  Davies, Walsh, \\& Booth~(1976) established that the observed diameter of \\sgra, the compact source in the Galactic center, scales as $\\lambda^2$, as expected if interstellar scattering from microstructure in the electron density determines the observed diameter.  The observed diameter of \\sgra\\ is now known to scale as $\\lambda^2$ from 30~cm to 3~mm (\\cite{rogersetal94}) and to be anisotropic at least over the wavelength range 21~cm to 7~mm (\\cite{bzkrml93}; \\cite{krichbaumetal93}; \\cite{y-zcwmr94}).  Maser spots in OH/IR stars within 25\\arcmin\\ of \\sgra\\ also show enhanced, anisotropic angular broadening (\\cite{vfcd92}; \\cite{fdcv94}).  These observations indicate that a region of enhanced scattering with an angular extent of at least 25\\arcmin\\ in radius (60~pc at 8.5~kpc) is along the line of sight to \\sgra.  At 1~GHz the level of angular broadening produced by this scattering region is roughly 10 times greater than that predicted by a recent model for the distribution of free electrons in the Galaxy (\\cite{tc93}, hereinafter TC93), even though this model includes a general enhancement of scattering toward the inner Galaxy.  These observations do not constrain the \\emph{radial} location of the scattering region for the following reason: All previous observations have been of sources in or near the Galactic center, and for such sources, a region of moderate scattering located far from the Galactic center can produce angular broadening equivalent to that from a region of intense scattering located close to the Galactic center.  Previous estimates for the location of the scattering region have ranged from~10~pc to~3~kpc.  Ozernoi \\& Shisov~(1977) concluded that an ``unrealistic'' level of turbulence is implied unless the region is within 10~pc of the Galactic center.  The level of turbulence they considered unrealistic, however, namely $\\sqrt{\\langle n_{\\mathrm{e}}^2\\rangle}/\\langle n_{\\mathrm{e}}\\rangle \\sim 1$, does appear to occur elsewhere in the interstellar medium (\\cite{s91}). Further, van~Langevelde et al.~(1992) used the free-free absorption toward \\sgra\\ to constrain the region's distance from the Galatic center to the range 0.85--3~kpc, though suitable adjustment of free parameters (outer scale and electron temperature) can decrease the limit to 0.03~kpc.  We shall refer to the case in which the region is a site of extreme scattering, $\\lesssim 100$~pc from the Galactic center and presumably caused by processes occurring there, as the GC model.  We shall refer to the case in which the region is far from the GC, $\\gtrsim 1$~kpc and a site of enhanced but not extreme scattering, as the random superposition (RS) model.  Although the GC model is attractive for phenomenological reasons, other sites of enhanced interstellar scattering are found throughout the Galaxy (e.g., NGC~6634, \\cite{mrgb90}; Cyg~X-3, \\cite{mmrj95}) and the mean free path for encountering such a region is approximately 8~kpc (\\cite{cwfsr91}).  Identifying the location of the scattering is important in establishing the origin of the scattering.  Associating the scattering with a specific region may elucidate the mechanism for the generation of the density fluctuations responsible for the scattering.  The currently favored mechanism is that velocity or magnetic field fluctuations---or both---generate the density fluctuations (\\cite{h84}, 1986; Montgomery, Brown, \\& Matthaeus~1987; \\cite{s91}; \\cite{sg94}; \\cite{gs95}).  Velocity or magnetic field fluctuations are also a natural means for inducing anisotropy in the density fluctuations and thereby in the scattering disks.  If this mechanism is correct, the amplitude of the density fluctuations may provide a measure of the coupling between the density and velocity or magnetic field fluctuations or, more generally, provide information about the small-scale velocity or magnetic field in the scattering region. However, current observational constraints are uncertain by the ratio of the Galactic center-scattering region distance to the Galactic center-Sun distance.  In the RS model, the ratio is a few while in the GC model the ratio could be as large as one hundred, so the location of the scattering region is a key free parameter.  The location of the scattering region also has implications for pulsar searches toward the \\hbox{GC}.  Cordes \\& Lazio~(1997) showed that even if the RS model is correct, pulsars seen through the scattering region will suffer pulse broadening of at least 5~s at 1~GHz (see also \\cite{dwb76}; \\cite{os77}).  If the GC model is correct, only at frequencies above 10~GHz will pulsations be detectable (because of the $\\nu^{-4}$ dependence of pulse broadening) and then only for pulsars with periods longer than 100~ms.  In this paper we develop a likelihood analysis to quantify the most probable $\\delgc$ for the scattering region.  In \\S\\ref{sec:gc.model} we describe our model for the distribution of free electrons in the \\hbox{GC}.  In \\S\\ref{sec:likefunc} we assemble measurements from the literature relevant to radio-wave scattering and develop a likelihood method to constrain the properties of the scattering region, and in \\S\\ref{sec:gc.conclude} we discuss our results and present our conclusions. ",
        "conclusions": "\\label{sec:gc.conclude}  \\subsection{Comparison with Previous Analyses}  Isaacman~(1981) surveyed the central $2\\arcdeg \\times 4\\arcdeg$ ($\\ell \\times b$) of the GC in a search for planetary nebulae.  He finds an \\emph{excess} number of sources as compared to that expected from extragalactic source counts, an excess he attributes to \\ion{H}{2} regions and planetary nebulae.  That he finds an excess at all is notable, though, since our analysis predicts that angular diameters of extragalactic sources seen through the GC scattering region will be at least 1\\farcm5--3\\arcmin.  The resolution of his survey was $0\\farcm4 \\times 2\\arcmin$, and he was able to detect sources with angular scales as large as 14\\arcmin.  Thus, we attribute his excess to the fact that, where his survey overlapped the GC scattering region, it was desensitized by the intense GC scattering to a considerably lesser degree than our survey.  Anantharamaiah et al.~(1991) used their observations at~0.327~GHz and a 0.408~GHz $\\log N$-$\\log S$ relation to conclude that the number of observed extragalactic sources within a 4~deg${}^2$ area centered on \\sgra\\ is consistent with the number expected from high-latitude source counts.  Outside of the inner 1~deg${}^2$, our source counts are also consistent with the expected number of extragalactic sources. Although scattering of other sources, such as B1739$-$298 and B1741$-$312 is heavy, the predicted diameter of these sources is less than 10\\arcsec\\ at~0.327~GHz, comparable to the size of their beam, so that these sources would not have been resolved out.  Gray et al.~(1993) surveyed the Sgr~E region ($\\ell = 358.7\\arcdeg, b = 0\\arcdeg$) at~0.843, 1.45, and~4.86~GHz.  At 1.4~GHz the number of sources they find is consistent with that expected from the $\\log N$-$\\log S$ distribution.  Figure~\\ref{fig:sclike} shows that the likelihood function for our field 358.9$+$0.5 does not favor a large amount of scattering, i.e., the number of sources in this field is consistent with that expected.  Further, two of the sources observed in our VLBI experiment (\\cite{lc97}), B1739$-$298 and 1LC~358.439$-$0.211, are in this field.  The former is heavily scattered, though not at a level sufficient for it to be seen through the \\sgra\\ scattering screen.  We conclude that our source count results are in good agreement with previous source counts toward the \\hbox{GC}.  The only exception occurs over the 1~deg${}^2$ region centered on \\sgra.  This region has not been considered previously or has been subsumed into a much larger area.  \\subsection{Physical Conditions in the Scattering Region}\\label{sec:gc.physical}  Our global likelihood, Fig.~\\ref{fig:global} and \\S\\ref{sec:gc.global}, attained a maximum for the following parameter values: $\\delgc = 150$~pc, $0.5\\arcdeg \\le \\psi_\\ell \\lesssim 1\\arcdeg$, and $l_0^{2/3}\\te^{-1/2} = 10^{-7}$.  In this section we consider whether a medium exists within the GC for which such parameter values are plausible.  There is a wealth of observational data available for the \\hbox{GC}.  We shall summarize those conclusions relevant to our study here; interested readers are referred to a number of recent reviews---Genzel, Hollenbach, \\& Townes~(1994); Morris \\& Serabyn~(1996); and Gredel~(1996)---and references within.  Our criteria for the host medium of the density fluctuations are that the medium must have a sufficient density and that it must be capable of sustaining density fluctuations of the requisite magnitude.  We establish our first criterion by estimating $\\nbar$ from the scattering diameters of GC sources using equations~(\\ref{eqn:measures}) and~(\\ref{eqn:smweight}).  The diameters of \\sgra\\ and the OH masers require a weighted scattering measure of $S \\approx 10^2$~\\smu, equation~(\\ref{eqn:smweight}). Eliminating SM between equations~(\\ref{eqn:measures}) and~(\\ref{eqn:smweight}) and solving for $\\nbar$ yields \\begin{equation} \\nbar \\sim 10^3\\,\\mathrm{cm}^{-3}\\frac{1}{\\varepsilon\\sqrt{f}}\\left(\\frac{l_0}{1\\,\\mathrm{pc}}\\right)^{1/3}\\left(\\frac{\\delgc}{150\\,\\mathrm{pc}}\\right)^{-3/2}. \\label{eqn:delne} \\end{equation}  Two factors could alter this estimate by about an order of magnitude. First, it is likely that $l_0 \\ll 1$~pc, which would \\emph{reduce} our estimate of $\\nbar$.  Second, as we noted in \\S\\ref{sec:gc.global}, the similarity between the values of $\\delgc$ and $\\psi_\\ell\\dgc$ suggests that the density fluctuations fill the region and $f \\approx 1$.  However, we might also associate $l_0$ with the characteristic size of a scattering cloudlet within the region.  If the region contains few such cloudlets and $\\delgc/l_0 \\gg 1$, then $f \\ll 1$, and our estimate above would be a considerable \\emph{underestimate}. Yusef-Zadeh et al.~(1994) estimated that a typical line of sight might intersect only 10 or so scattering cloudlets.  In any event, we conclude that the scattering medium must be dense, $n_{\\mathrm{e}} \\gtrsim 10^2$~cm${}^{-3}$.  For comparison, Spangler~(1991) concludes that $n_{\\mathrm{e}} \\sim 1$~cm${}^{-3}$ for scattering regions in the Galactic disk.  Our second criterion for the host medium is that it must be able to support density fluctuations of the required magnitude.  This constraint has been lucidly reviewed by Spangler~(1991): The density fluctuations are presumed to arise from plasma turbulence.  As this turbulence dissipates, it cannot heat the host medium at a rate that exceeds the medium's cooling capacity.  This constraint is particularly acute in the situation we are proposing as the dissipation mechanisms considered by Spangler~(1991) scale as $l_0^{-a}$ with $a \\approx 1$.  Since we are considering $l_0 < 1$~pc, the heating rates could be excessive.  The dominant damping mechanisms for $l_0 < 1$~pc are linear Landau damping, ion-neutral collisions, and a parametric decay instability. The first two mechanisms scale as $l_0^{-2/3}$ while the latter scales as $l_0^{-1}$.  In addition to their dependence on $l_0$, the damping rates depend on the large scale magnetic field, $\\Gamma \\propto B^2$; the Alfv\\'en wave speed, $\\va$; and the amplitude of the magnetic fluctuations, $\\Gamma \\propto (\\delta B/B)^2$ for linear Landau damping and ion-neutral collisions while $\\Gamma \\propto (\\delta B/B)^3$ for the parametric decay instability.  Linear Landau damping also depends upon the angle of propagation with respect to the direction of $\\mathbf{B}$, $\\chi$, and the plasma $\\beta$.  For values of these quantities appropriate for scattering regions in the Galactic disk, these damping mechanisms produce volumetric heating rates of $\\Gamma \\sim 10^{-23.5}$--$10^{-21.5}$~erg~s${}^{-1}$~cm${}^{-3}$.  As Spangler~(1991) discussed, there are also a number of simplifications and additional assumptions which enter the calculation of these heating rates.  Inferred magnetic field strengths in the GC are $B \\sim 1$~mG, or $10^3$ that of the field strength in the disk.  To estimate $\\delta B/B$, we use (Cordes, Clegg, \\& Simonetti~1990) \\begin{equation} \\frac{\\delne}{n_{\\mathrm{e}}} \\sim \\left(\\frac{\\delta B}{B}\\right)^c \\label{eqn:bB} \\end{equation} with $c = 1$ for linear processes and $c = 2$ for non-linear processes like the parametric decay instability.  Our first criterion for the host medium is that $\\delne/n_{\\mathrm{e}} \\le 1$.  For definiteness, and to provide the largest possible value of the heating, we take $\\delta B/B \\sim 1$.  Finally, although $B$ is much larger in the GC than in the Galactic disk, $n$ is also larger.  As a result $\\va$ is larger than in the disk, but probably by no more than an order of magnitude.  Thus, we expect $\\Gamma$ in the GC to be about a factor of $10^7$ larger that in the Galactic disk.  The heating rate from linear Landau damping in scattering regions in the Galactic disk (\\cite{s91}) assumes the density fluctuations arise from obliquely propagating magnetosonic waves ($\\chi \\approx 6\\arcdeg$).  More aligned propagation results in less damping.  It is not clear if the GC environment would favor highly aligned propagation or not.  The large values of $B$ in the GC are inferred, in part, from the system of non-thermal filaments and threads seen throughout the \\hbox{GC}.  With only one exception, these filaments have no kinks or bends in them, even though they are observed to be interacting with molecular clouds having typical velocities of 10--100~km~s${}^{-1}$ (\\cite{ms96}).  This rigidity could be an indication that only highly aligned propagation is allowed.  If $\\chi$ is highly concentrated near 0\\arcdeg, then the heating from linear Landau damping would be unimportant and the heating rates could be two orders of magnitude lower than those quoted above.  Alternately, as Spangler~(1991) noted, the distribution of $\\chi$ could be isotropic, but waves with large $\\chi$ would then damp quickly and the heating rate will be unchanged or even larger than what we assume.  The presence of small-scale ($\\approx 0.1$~pc) magnetoionic cloudlets in the GC has already been inferred to explain large changes in the Faraday rotation measure of certain features (G~359.1$-$00.2, the ``Snake,'' \\cite{gnec95}; G~359.54$+$0.18, the non-thermal filaments, Yusef-Zadeh, Wardle, \\& Parastaran~1997), though the inferred density in these cloudlets, 0.3--10~cm${}^{-3}$, is less than our nominal estimate.  In the Galactic disk a small body of observational evidence suggests that the magnetoionic medium responsible for Faraday rotation is also responsible for scattering and pulsar dispersion (\\cite{sc86}; Lazio, Spangler, \\& Cordes~1990; \\cite{ahmsrs96}), and the same may be true in the \\hbox{GC}.  We now consider two models for the host medium.  In both models, the scattering arises in thin layers on the surfaces of molecular clouds. Even if the filling factor, $f$, of these layers is not large, the \\emph{covering factor}, i.e., the probability that a line of sight through the GC will intersect one of these layers, can still be close to unity.  \\subsubsection{Photoionized Surfaces of Molecular Clouds}\\label{sec:gc.skins}  Over the region $|\\ell| \\lesssim 1\\fdg5$ and $|b| \\lesssim 0\\fdg5$, $n_{\\mathrm{e}} \\sim 10$~cm${}^{-3}$, as determined from single-dish recombination line and total intensity measurements (Matthews, Davies, \\& Pedlar~1973; \\cite{mp79}).  Embedded within this large-scale region are smaller regions of much higher electron densities, $n_{\\mathrm{e}} \\sim 10^3$--$10^5$~cm${}^{-3}$, primarily within Sgr~A West and Sgr~B2 (e.g., \\cite{mpgy-z93}).  Some of these high density regions are the photoionized surfaces (size $\\sim 10^{-4}$~pc) of molecular clouds ($n \\gtrsim 10^4$~cm${}^{-3}$) irradiated by the ambient radiation field (effective temperature $\\approx 35\\,000$~K).  Yusef-Zadeh et al.~(1994) identified these molecular skins as the source of the scattering and associated their thicknesses with the outer scale, $l_0$; Gray et al.~(1995) suggested that the magnetoionic medium responsible for the Faraday rotation toward G~359.1$-$00.2 also results from these molecular clouds.  This model suffers from at least three potential difficulties.  First, the molecular skins are photoionized by a radiation field having a temperature of $T_{\\mathrm{eff}} \\sim 10^4$~\\hbox{K}.  Our constraint on $l_0^{2/3}\\te^{-1/2}$ therefore requires $l_0 \\sim 10^{-7.1}$~pc. This value is considerably smaller than that derived by Yusef-Zadeh et al.~(1994) for the ionized molecular skins.  However, in deriving their value for $l_0$, Yusef-Zadeh et al.~(1994) used a value of the ionizing flux appropriate to the inner few parsecs.  The stellar density decreases as $r^{-2}$, so outside the inner few parsecs, the ionizing flux should be lower than that assumed by Yusef-Zadeh et al.~(1994).  A lower ionizing flux would result in a smaller skin depth and bring their estimate and our estimate of $l_0$ into better agreement.  A second potential difficulty with this model is that the medium would only barely be capable of cooling itself.  If the outer scale is $l_0 \\sim 10^{-7}$~pc, then the heating rate from the damping of the plasma turbulence is $\\Gamma \\sim 10^{-13}$--$10^{-12}$~erg~cm${}^{-3}$~s${}^{-1}$.  The cooling capacity of gas near $\\te \\sim 10^4$~K depends sensitively upon the fractional ionization and temperature (\\cite{dm72}).  We estimate that a density of $n_{\\mathrm{e}} \\gtrsim 10^5$~cm${}^{-3}$ is required for these skins to be able to cool sufficiently in order to support the density fluctuations.  This density is at the upper end of the range $10^3$--$10^5$~cm${}^{-3}$ inferred for the small-scale \\ion{H}{2} regions.  In determining the cooling function of the medium, we have used results that assume a solar abundance.  The metallicity in the GC could be as much as twice solar, leading to an increased cooling efficiency.  The third difficulty is that the required value for $l_0$ in this model, $l_0 \\sim 3 \\times 10^{11}$~cm, is considerably smaller than that in the Galactic disk.  In the Galactic disk, a stringent lower limit on the outer scale is $10^{13}$~cm, and it may be as large as $10^{18}$~cm (\\cite{ars95}).  Although the physics for the generation and maintenance of small-scale density fluctuations is not well understood, we regard it as potentially troublesome that this model predicts such a small $l_0$.  \\subsubsection{``Warm'' Interfaces}\\label{sec:gc.interface}  X-ray observations have revealed a central, diffuse X-ray source with a (FWHM) size of $1\\fdg8 \\times 0\\fdg9$ (\\cite{ykkkt90}).  Frail et al.~(1994) suggested that this X-ray emitting gas may be responsible for the scattering.  Yusef-Zadeh et al.~(1997) suggested that this X-ray emitting gas is also the magnetoionic medium responsible for the Faraday rotation toward G~359.54$+$0.18.  The density and temperature of this region are estimated at 0.05~cm${}^{-3}$ and $10^7$--$10^8$~\\hbox{K}.  This region cannot itself be the host of the density fluctuations because of its low density, cf.~equation~(\\ref{eqn:delne}).  However, this gas appears spatially coincident with the central zone of intense molecular emission and presumably abuts cooler gas in the clouds.  We modify Frail et al.'s~(1994) proposal by identifying the interfaces where the GC molecular clouds are exposed to this ambient hot medium as the source of the scattering.  We term these interfaces ``warm'' by analogy with McKee \\& Ostriker's~(1977) model for the \\hbox{ISM}.  In that model cold clouds immersed in a hot ($10^6$~K) medium have $10^4$~K interfaces.  In the GC densities and temperatures are 1--2 orders of magnitude higher, but we expect that clouds will still develop intermediate temperature interfaces.  This model suffers from two of the same difficulties as the previous model.  The X-ray emitting gas and molecular clouds appear to be in rough pressure equilibrium with $P \\sim 5 \\times 10^6$~K~cm${}^{-3}$ (\\cite{bs91}).  Even though there are supersonic motions \\emph{within} the clouds, the clouds themselves ($v \\sim 10$--100~km~s${}^{-1}$) are moving subsonically with respect to the hot medium ($c_{\\mathrm{sound}} \\sim 1000$~km~s${}^{-1}$).  Taking pressure balance to extend throughout the interface region, we find a density $n_{\\mathrm{e}} \\sim 5$--50~cm${}^{-3}$ for $\\te \\sim 10^5$--$10^6$~\\hbox{K}.  From our likelihood results, the temperature within these interfaces implies an outer scale of $l_0 \\sim 10^{-6.5}$--$10^{-6}$~pc, which, in turn, implies an rms density, equation~(\\ref{eqn:delne}), of $\\delne \\sim 10$~cm${}^{-3}$.  However, the cooling capacity of this medium is only $10^{-20}$~erg~cm${}^{-3}$~s${}^{-1}$.  The predicted heating rate is $\\Gamma \\sim 10^{-13}$~erg~cm${}^{-3}$~s${}^{-1}$.  The outer scale in this model remains troublesomely small.  If the size of the interface region is set by thermal conduction, the portion of the interface with $\\te < 10^6~K$ has a size $\\lesssim 10^{-1}$~pc (\\cite{mc77}).  Clearly if not all of the interface contributes to the scattering better agreement would be obtained between $l_0$ and the interface size.  Still, the outer scale remains an order of magnitude smaller than its lower limit in the Galactic disk.  One point in favor of this model is the distribution of the X-ray emitting gas as compared to that of the molecular clouds.  The size of the X-ray emitting region is similar to the extent of the scattering region, approximately 1\\arcdeg.  In contrast, the molecular cloud distribution extends over the range $-1\\arcdeg \\lesssim \\ell \\lesssim 2\\arcdeg$.  If the scattering traced massive stars within these molecular clouds, as the photoionized molecular cloud skins model suggests, the scattering should extend further in longitude than it does.  In this respect, the lack of enhanced scattering for \\hoh\\ masers in Sgr~B is particularly problematic.  We stress the importance, and probable uniqueness, of the high density in the \\hbox{GC}.  In the Galactic disk density fluctuations cannot be supported in media with $\\te \\sim 10^6$~K, a position with both theoretical and limited observational support (\\cite{s91}; \\cite{pc92}).  Similarly, recent VLBI observations of 5 pulsars show no evidence for an enhanced level of turbulence at the boundary of the Local Bubble (\\cite{cr87}; Britton, Gwinn, \\& Ojeda~1996), potentially a local analog of an interface between a hot and cooler medium. However, the Local Bubble and ambient medium have densities a factor of $10^2$--$10^3$ smaller than that in the \\hbox{GC}.  In summary, we use our likelihood results, \\S\\ref{sec:gc.global}, to constrain host media for the scattering material.  Potential media include the photoionized skins of molecular clouds or the interface regions between the clouds and the ambient X-ray emitting gas.  There are difficulties with both models: Both models overpredict the outer scale and appear to have some trouble supporting the required level of density fluctuations.  Although we have been unable to make an unambiguous identification of the scattering medium with either medium, we favor the interface model, in part, because it shows a better correspondence between the spatial distribution of scattering and proposed host medium.  \\subsection{Modification of the Taylor-Cordes Model}\\label{sec:tcmodel}  The TC93 model modelled the global distribution of free electrons in the Galaxy with four components: an extended component, an inner Galaxy component, a component confined to the spiral arms, and a component local to the Gum Nebula.  We now extend the TC93 model to include a GC component\\footnote{ As in TC93, the coordinate system has the $x$-axis directed parallel to $\\ell = 90\\arcdeg$, the $y$-axis toward $\\ell = 180\\arcdeg$, the $z$-axis toward $b = 90\\arcdeg$, and $R = \\sqrt{x^2 + y^2}$ is the Galactocentric radius.} (cf.\\ eqn.~[11] of TC93): \\begin{eqnarray} n_{\\mathrm{e}}(x,y,z) & = & n_1(R,z) + n_2(R, z) + \\sum_{j=1}^4 n_{\\mathrm{arm},j}(x,y,z) + n_{\\mathrm{Gum}}(x,y,z) \\nonumber \\\\ & + & n_{\\mathrm{GC}}g_{\\mathrm{GC}}(R)h_{\\mathrm{GC}}(z). \\end{eqnarray} The first four components are discussed at length in TC93.  We focus on only the last component, that toward the \\hbox{GC}.  Based on the estimate in equation~(\\ref{eqn:delne}) and our estimates for $l_0$, we take $n_{\\mathrm{GC}} = 10$~cm${}^{-3}$.  Our estimate for $\\delgc$ is $\\delgc = 150$~pc.  Heretofore, we have been treating the scattering region as a screen with sharp boundaries.  It is more likely that the region has soft edges.  We therefore adopt a radial dependence of \\begin{equation} g_{\\mathrm{GC}}(R) = \\exp\\left[-(R/0.150\\,\\mathrm{kpc})^2\\right]. \\end{equation}  The latitude (or $z$) dependence of the screen is less well constrained.  For definiteness we take \\begin{equation} h_{\\mathrm{GC}}(z) = \\exp\\left[-(z/0.075\\,\\mathrm{kpc})^2\\right] \\end{equation} corresponding to $\\psi_b = 0\\fdg5$.  The resulting axial ratio for the electron density distribution is 0.5; the axial ratio for the X-ray distribution is also 0.5 and that of the molecular cloud distribution is 0.3 (\\cite{ms96}).  In the TC93 model the relationship between the free electron density and the scattering measure produced by a line of sight of length $ds$ through those electrons is $d\\mathrm{SM} \\propto F n_{\\mathrm{e}}^2ds$.  The parameter~$F$ is \\begin{equation} F = \\frac{\\zeta\\epsilon^2}{f}\\left(\\frac{l_0}{1\\,\\mathrm{pc}}\\right)^2, \\label{eqn:F} \\end{equation} where $\\zeta$ is the normalized variance of electron density fluctuations between cloudlets and their surroundings and $\\epsilon$ and~$f$ are as in equation~(\\ref{eqn:measures}).  Taking $\\zeta \\sim \\epsilon \\sim 1$, our estimates for $l_0$ imply $F \\gtrsim 10^4$. Both of the models we have considered here have $f < 1$.  For definiteness, we take $f \\sim 0.1$, recognizing that this may be an upper limit on~$f$.  We therefore conclude $F \\gtrsim 10^5$.  For comparison, the parameter~$F$ has a value of~0.4 in the solar neighborhood, 6 in spiral arms, and 40 in the Galaxy's inner few kiloparsecs.  A value of $F \\sim 10^5$ produces an SM comparable to that suggested by the scattering diameters of \\sgra\\ and the OH masers.  Their scattering diameters require a line-of-sight weighted scattering measure of $S \\approx 10^2$~kpc~m${}^{-20/3}$.  Our results suggest $\\delgc \\approx 150$~pc; correcting for the line-of-sight weighting, equation~(\\ref{eqn:smweight}), implies that the GC has a scattering measure of $\\mathrm{SM} \\sim 10^{5.5}$~kpc~m${}^{-20/3}$.  Integrating the TC93 expression for $d\\mathrm{SM}$ through the GC, with $F \\sim 10^5$, we find $\\mathrm{SM} \\sim 10^6$~kpc~m${}^{-20/3}$.  Since the GC component is so localized, only for lines of sight through the GC are the results of TC93 altered.  In this direction, however, the TC93 model underpredicts various quantities by a large amount.  In the TC93 model, GC pulsars have $\\mathrm{DM} \\approx 600$--800~pc~cm${}^{-3}$; we predict that the DM will be somewhat larger, $\\mathrm{DM} \\approx 2000$~pc~cm${}^{-3}$, with approximately 1500~pc~cm${}^{-3}$ of that arising from the GC component itself.  For comparison, the largest DM known is for PSR~B1758$-$23 with $\\mathrm{DM} = 1074$~pc~cm${}^{-3}$ (Manchester, D'Amico, \\& Tuohy~1985; \\cite{klmjds93}).  Further, from Cordes \\& Lazio~(1997) the temporal broadening of pulses from pulsars seen through this region will be $350\\,\\nu_{\\mathrm{GHz}}^{-4}$~\\emph{seconds}, requiring high-frequency ($\\nu \\approx 10$~GHz) periodicity searches to detect pulsations.  Although the DM we predict for GC pulsars is substantial, the dispersion smearing across a 1~GHz bandpass at 10~GHz ($\\approx 5$~ms) is comparable to the pulse broadening, so that only a small number of filterbank channels, e.g., 16, would be necessary to combat the dispersion smearing.  \\subsection{Conclusions}  We use a likelihood analysis to determine the following parameters of the GC scattering region: The GC-scattering region separation, $\\delgc$; the angular extent of the region, $\\psi_\\ell$ and $\\psi_b$; the outer scale on which density fluctuations occur, $l_0$; and the gas temperature, $\\te$. \\begin{itemize} \\item From the literature we have assembled a list of all sources toward the GC for which angular broadening has been measured.  A subset of these sources is OH/IR stars, for which the spatial distribution about the GC is known.  We construct a likelihood function for the angular broadening of OH/IR stars, utilizing this distribution, \\S\\ref{sec:broaden}.  Masers within approximately 1\\arcdeg\\ of \\sgra\\ have diameters consistent with $\\delgc \\approx 150$~pc (Fig.~\\ref{fig:aggregate_as}).  \\item The likelihood analysis of our source counts, \\S\\ref{sec:counts}, indicates that a deficit of sources occurs within approximately 1\\arcdeg\\ of \\sgra\\ and is caused by a scattering region within 500~pc of \\sgra\\ (Fig.~\\ref{fig:aggregate_sc}).  The resulting scattering diameter, at least 20\\arcsec\\ at~1~GHz, causes extragalactic sources to be so broad as to be resolved out by our observations.  \\item \\hoh\\ masers in and an extragalactic source near Sgr~B and an extragalactic source near Sgr~C do not show the extreme scattering of sources closer to \\sgra, indicating that the scattering region does not extend to more than 1\\arcdeg\\ in longitude.  The latitude extent of the scattering region is poorly constrained, but is no more than 1\\arcdeg.  \\item From the literature we have estimated the free-free emission and absorption toward five heavily scattered masers near \\sgra.  The likelihood function is dominated by the free-free emission.  The relevant parameters, $\\delgc$, $l_0$, and $\\te$, are not independent for this likelihood and only the product $\\delgc^{-2}l_0^{2/3}\\te^{-1/2}$ can be constrained (Fig.~\\ref{fig:aggregate_ff}). \\end{itemize}  The global likelihood, formed by multiplying the individual likelihoods, is shown in Fig.~\\ref{fig:global}.  The maximum likelihood estimates of the parameters are $\\delgc = 150$~pc, $0.5\\arcdeg \\le \\psi_\\ell \\lesssim 1\\arcdeg$, and $l_0^{2/3}\\te^{-1/2} = 10^{-7}$ with $l_0$ in pc and $\\te$ in \\hbox{K}.  The parameter $\\psi_b$ was not well constrained and we adopt $\\psi_b = 0\\fdg5$.  The close correspondence between $\\delgc$ and $\\psi_\\ell\\dgc$ suggests that the scattering region encloses the \\hbox{GC}.  The GC scattering region produces a 1~GHz scattering diameter of 90\\arcsec, if the region is a single screen, or 180\\arcsec, if the region wraps around the GC, as appears probable.  We modify the Taylor-Cordes model for the Galactic distribution of free electrons in order to include an explicit GC component.  We predict that pulsars seen through this region will have a dispersion measure of approximately $2000$~pc~cm${}^{-3}$, of which approximately 1500~pc~cm${}^{-3}$ arises from the GC component itself, and suffer pulse broadening of $350\\,\\nu_{\\mathrm{GHz}}^{-4}$~\\emph{seconds}; pulsations will be detected only for frequencies above 10~GHz (\\cite{cl97}).  As host media for the scattering we consider the photoionized surface layers of molecular clouds and the interfaces between molecular clouds and the $10^7$~K ambient gas.  We identify the host medium by requiring that it be sufficiently dense to support density fluctuations of the required magnitude.  We are unable to make an unambiguous determination, but we favor the interface model which predicts that the scattering medium is hot ($\\te \\sim 10^6$~K) and dense ($n_{\\mathrm{e}} \\sim 10$~cm${}^{-3}$).  The X-ray interface model also shows better spatial agreement, when compared to the photoionized skin model, with the region over which the scattering is observed.  This model is summarized graphically in Fig.~\\ref{fig:gc.summary}.  The GC scattering region is likely to be unique in the Galaxy, probably because it is a high-pressure environment and can sustain densities and temperatures much higher than in the Galactic disk."
    },
    "9804/astro-ph9804294_arXiv.txt": {
        "abstract": "{ Lithium is one of the few primordially produced elements. The value of the primordial Li is taken to be that observed in metal--poor dwarfs, where  it is  not  contaminated by stellar Li sources which act on longer time scales. The atmospheric abundance is currently derived  from the LiI $\\lambda\\lambda 6707 \\AA~$  resonance transition and the validity of the models employed has been questioned \\markcite{k95} (Kurucz 1995). In this letter we report the first detection of the Li I $\\lambda\\lambda   6104 \\AA~ 2^2P - 3^2D$  subordinate transition in the prototype  population II star HD~140283. The same Li abundance of (Li/H) $=1.4\\times 10^{-10}$ is found consistent with both the resonance and subordinate lines. The two lines    form at different depths in the atmosphere implying  that the 1-D homogeneous atmospheric models used in the abundance determination are essentially correct. When coupled with the standard big bang yields, the  Li in the halo dwarfs provides two solutions for the baryon-to-photon ratio $\\eta_{10}= n_{b}/n_{\\gamma} \\times 10^{10}$ and for the present baryon density $\\Omega_b h_{70}^2=0.0748\\eta_{10}$: a) a first solution  at  $\\eta_{10}\\approx 1.8$, consistent with   the $\\eta_{10}$ implied by the high deuterium values $D/H\\approx 2\\times 10^{-4}$ observed in some quasar absorption systems \\markcite{webb} (Webb et al 1997) and b) a second solution  at $\\eta_{10}$ $\\approx$  4 which is  consistent, within the errors, with    the low  deuterium D/H =$3.4\\times 10^{-5}$ measured in other quasar absorption systems\\markcite{burles} (Burles \\& Tytler 1998). } ",
        "introduction": "Lithium,  together with D and $\\rm ^{3,4}He$, is one of the few elements produced by nuclear reactions in the first minutes after big bang\\markcite{wfh} (Wagoner, Fowler \\& Hoyle 1967). The observations of these elements and their extrapolation to the primordial values are consistent with the predictions of the standard primordial nucleosynthesis providing, together with the relic radiation and the expansion of the Universe, a robust support to the big bang theory. Recently, additional support to the primordial nature of Li in halo dwarfs has come from the observations of Li in metal-poor stars of the thick disc\\markcite{mbp97} (Molaro, Bonifacio \\& Pasquini 1997). This population is chemically and kinematically distinct from the halo, but has  the same Li abundance of the halo. Minniti et al (1997) \\markcite{min97} claimed detection of Li, at the plateau level, in a metal-rich, but old star, belonging to the Galactic Bulge. Finally Li at the plateau level has also been detected in a star which was possibly born in an external galaxy and then accreted   by the Milky Way\\markcite{mol97} (Molaro 1997).  So far the Li abundance has been  always obtained only from the analysis of the Li I $\\lambda\\lambda$ 6707 \\AA~ resonance doublet. This is not a very comfortable situation in the light of the importance of the  determination of lithium  abundances in stars for primordial nucleosynthesis,  stellar structure and chemical evolution.  Our ability to determine the Li abundance using simple plane-parallel homogeneous atmospheres, has been recently debated\\markcite{k95,kis97,gp97} (Kurucz 1995; Kiselman 1997; Gadun \\& Pavlenko 1997). The analysis of several  lines, which sample different depths in the stellar atmosphere is crucial to test the correctness of the modelling. The one dimensional, homogeneous, static models which are currently employed may arise concern because they ignore the fine structure and hydrodynamic phenomena such as granulation which are seen on the Sun.  The Li I $\\lambda\\lambda$ 6707 \\AA ~resonance transition is the only one readily available to spectroscopic observation. The strongest subordinate line at 6104 \\AA ~ is  much  fainter and blended with Fe I line and has been so far detected only in young T Tauri stars (Hartigan et al 1989)\\markcite{hart} and Li-rich giants (Merchant 1967, Wallerstein \\& Sneden 1982) \\markcite{merchant,ws}, were  Li is more than about 1 dex more abundant owing to the Galactic Li production. \\par In this letter we report the detection of the Li I $\\lambda\\lambda$ 6104 \\AA ~transition in the spectrum of the metal--poor star HD140283. Both this line and the resonance line are consistent with the computations made using a one dimensional, homogeneous model atmosphere, thus increasing our confidence that this model represents a satisfactory average of the complex fine structure expected in metal--poor stars.  The use of Li observed in halo dwarfs as an indicator of primordial abundance rests on the absence of any Li  depletion. Depletion is predicted by non-standard models which take into account rotational mixing \\markcite{pin92} (Pinsonneault, Deliyannis \\& Demarque 1992) or diffusion \\markcite{vau95} (Vauclair \\& Charbonnel 1995), but these models predict a downturn of the hot side of the Li plateau and considerable dispersion. It seems that neither the downturn nor the large dispersion is present in the observations, which suggests that diffusion or rotational mixing do not affect significantly the Li observed at the stellar surface of metal--poor dwarfs. However, the downturn can be very small ($\\approx 0.2$ dex) for the purely diffusive case and a suitable choice of the mixing length parameter ($\\alpha=1.5$, see Deliyannis et al 1990\\markcite{deli90}) and the issue of intrinsic dispersion remains rather controversial with some positive claims. Ryan et al (1996)\\markcite{ryan} identify a triplet of stars (G064-012, G064-037, CD -33$^\\circ$ 1173) with similar colors, but different Li abundances by a factor of 2.5. Then there is  the case of star BD+23 3912 which has a [Fe/H]$\\approx -1.3$ to  $-1.5$ and  a Li abundance which is about 0.20-0.36 dex higher than the plateau (Rebolo et al 1988, King et al 1996\\markcite{reb88,king96}). Moreover Boesgaard et al (1998)\\markcite{boesgaard} find differences of up to $\\approx 0.5$ dex among seven subgiants of M92 but the same objects show other chemical peculiarities, namely [Mg/Fe] is 0.55 dex lower and [Na/Fe] is 0.76 dex larger than in HD140283 (King et al 1998) \\markcite{king98} . ",
        "conclusions": "The LiI line forming regions lie  in the upper part of the atmospheric convective zone where Li is mostly ionized due to its low ionization potential. This is why the determination of precise Li abundances requires accurate observations, accurate stellar effective temperature and an appropriate modeling of the atmosphere of a metal poor  star. The model-atmospheres employed are one dimensional (1-D), with plane parallel geometry and ignore any inhomogeneity effect, such as granulation. Qualitative computations,  based on a two-stream model atmosphere, suggested that the abundance of Li in halo dwarfs could be underestimated by as much as a factor of 10\\markcite{k95} (Kurucz 1995), but more recent calculations based on 2-D\\markcite{gp97} (Gadun \\& Pavlenko 1997) and 3-D\\markcite{kis97} (Kiselman 1997) atmospheric models show that effects of granulation on the LiI lines are much less important. Granulation effects in the atmosphere have a depth dependence  and this should produce different effects in the resonance and subordinate doublets.  As  can be seen from  figures 1 and 2, the same Li abundance reproduces satisfactorily both the 6104 \\AA ~and the 6707 \\AA ~doublets. The two transitions form at different depths in the stellar atmosphere: unit optical depth at wavelength 6707.761 \\AA~ is attained at $log (\\tau_{Ross})\\approx -0.57$, corresponding in our model to a local temperature of 5235 K, while at wavelength 6103.649 \\AA~ it is already attained at $log (\\tau_{Ross})\\approx -0.09$, or T=5915 K. The resonance line receives contributions from a more extended region than the subordinate line. Unit optical depth at the  wavelength at which the residual intensity is 0.999, is attained at $log (\\tau_{Ross})\\approx -0.11$ for the resonance line, but at $log (\\tau_{Ross})\\approx -0.08$ for the subordinate line. Thus the subordinate line samples deeper and hotter layers than the resonance line, as shown in Fig. 3. \\par   The lower level of the Li 6104 \\AA~ transition is the upper level of the 6707 \\AA $2^2S-2^2P$~ transition. Our synthetic spectra are computed under the LTE assumption and the  consistency between the two lines implies a correct computation of  the populations of the 2S, 2P and 3D levels. This is in agreement with the theoretical estimations which predict relatively small corrections for NLTE effects in the LI 6707 \\AA ~line \\markcite{carl94,pavmag} (Carlsson et al 1994; Pavlenko \\& Magazz\\`u 1996). Thus the  detection of a subordinate LiI line, and its consistency with the abundance derived from the resonance 6707 \\AA~ doublet, provides support to the correctness of this Li abundance.  The consistency of the abundances based on the LiI 6707 \\AA~ and 6104 \\AA~  transitions observed in HD 140283 supports  the Li abundances measured in the population II stars, using 1-D model atmospheres, in the last decades. The new generation of large  telescopes will allow  to measure the Li 6104 \\AA ~Li I  subordinate doublet in other much fainter population II stars, thus permitting  to verify  this consistency on the grounds  of  a statistically  significant  sample, and ultimately achieve a more accurate measurement of the primordial Li abundance. \\par Among the light elements produced in the first minutes after the big bang, Li is the only one which shows a non monotonic behaviour with $\\eta_{10}$, the so-called  {\\sl Li-valley}, which reflects the different nuclear reactions which synthesize Li at different baryonic densities. The most recent measurement of the Li primordial abundance is $\\rm (Li/H)=1.73 \\pm 0.05_{stat} \\pm 0.2_{sys} \\times 10^{-10}$\\markcite{bm} (Bonifacio \\& Molaro 1997), which is the mean value of 41 halo stars for which precise effective temperatures, determined by means of the infrared flux method \\markcite{alonso} ({Alonso}, {Arribas}, \\& {Martinez-Roger}  1996), were available. The systematic errors, which dominate the error budget,   come from a possible offset of $\\pm$ 75 K in the zero point of the temperature scale of cool stars. This Li abundance intercepts the primordial yields for two different values of $\\eta_{10}$, which unfortunately do not help in resolving the deuterium and helium controversies. Each  solution for $\\eta_{10}$ obtained from Li is consistent with either the high-deuterium/low-helium\\markcite{webb,olive} (Webb et al 1997; Olive, Steigman \\& Skillman 1997) or the low-deuterium/high-helium \\markcite{burles,izo97} (Burles \\& Tytler 1998; Izotov, Thuan \\& Lipovetsky 1997). The lower $\\eta_{10}$ requires considerable D destruction to match the presently observed abundance in the local interstellar medium of the Galaxy. The higher $\\eta_{10}$ value is also consistent with the low deuterium (D/H $=3.9 (\\pm 1)\\times 10^{-5}$) derived from the 92 cm hyperfine transition emission towards the unprocessed Galactic anti-center \\markcite{chen97} (Chengalur, Braun \\& Butler Burton 1997)."
    },
    "9804/astro-ph9804249_arXiv.txt": {
        "abstract": "Modifications in Friedmann-Lema\\^itre-Robertson-Walker (FLRW) Hubble diagrams caused by  mass density inhomogeneities are used to illustrate possible effects on a determination of the mass parameter $\\OM$ and the cosmological constant $\\Lambda$. The values of these parameters inferred from a given set of observations depend on the fractional amount of matter in inhomogeneities and can  differ from those obtained by using standard FLRW Hubble diagrams by as much as a factor of two. ",
        "introduction": "The pressure-free FLRW models of GR have, for a very long time,  been assumed to adequately describe the large scale geometry of our universe but only recently have observational techniques  become available which promise to determine values for all three of their requisite parameters $\\{H_0,\\ \\OM$, and $\\OL\\equiv {\\Lambda c^2/(3H_0^2)}\\}$ and implicitly test this assumption.  As an example,  corrected magnitudes and redshifts ($m$-$z$) for Type Ia supernovae (SNe Ia) are measured, plotted, and compared with theoretical $m(H_0,\\ \\OM,\\OL; z)$ curves computed for the FLRW models \\cite{PS1,PS2,GP}. Because these models are isotropic and homogeneous, and our universe appears quite inhomogeneous, modifications in the FLRW predictions have long been proposed and estimated \\cite{YZ,KR}. When a wide angle measurement  of the CMB is made the average mass density of the universe (the FLRW value for $\\rho_0$) likely exists within the radio beam collected by the antenna. However, when small objects such as SNe Ia are observed at $z< 1$, a mass density significantly less than the average is `likely' to be in the observing optical path. In particular, if the underlying mass density approximately follows luminous matter (\\ie associated with bounded galaxies) then effects of a diminished mass density in the observing beams on relations like $m(\\OM,\\OL;z)$ are important.  The majority of currently observed SNe Ia are not being seen through foreground galaxies and whether or not this is due to selection (rather than statistics) is not important.  If the objects observed do not have the average FLRW mass density $\\rho_0$ in their foregrounds then the FLRW \\mz\\ relation does not apply to them. Ultimately some SNe Ia should exist behind foreground galaxies and for these, \\mz\\ should be computed using the lensing formulas.  These formulas \\cite{BR,CJ} contain source-observer, deflector-observer, and source-deflector distances, respectively $D_s, D_d$, and $D_{ds}$, all of which depend on the mass density in the observing beam, {\\bf excluding} the deflector. These distances will not be given by the standard FLRW result if the observing beam contains less than the average FLRW mass density, but instead given by the `intergalactic' distance discussed here. ",
        "conclusions": "Numerous arguments have been made since the '60s against the existence of any effect on apparent magnitudes such as given here (see \\cite{KR2} for current refs. to dissenting opinions). These are almost all weak-lensing arguments, valid as long as density perturbations don't produce multiple images (or absorb photons), \\ie the FLRW result will coincide with the theoretical mean of weak lensing observations.  Even if weak lensing arguments can be extended to strong lensing perturbations, \\eg by not resolving separate images, the theoretical mean of a distribution of magnitudes may be of little use in determining $\\OM$ or $\\OL$. If mass is as inhomogeneous as luminosity, the distribution of magnitudes is expected to be so skewed as to make  the mean statistically insignificant. The `most likely' value should be far more useful in a determination of the cosmic parameters. However, to determine the most likely \\mz not only requires knowledge of the average mass density $\\rho_0$, it requires the modeling of galaxies masses etc. What can be determined by fixing a single additional parameter $\\nu$ (which proportions total FLRW mass density into  an intergalactic  component and a galactic component)  is the `intergalactic' Hubble curve. This \\mz can be much more useful in determining $\\OM$ or $\\OL$ than the mean Hubble curve, \\eg if galaxies are not larger than they appear optically, the distribution of magnitudes (at a given redshift) is expected to peak much closer to the intergalactic value than to the mean value (see the numerical work \\cite{HD}). Additionally the effects of galaxy lensing can be controlled by simple selection. All that is required is that the observed SNe Ia are separated into those with foreground galaxies and those without. Those without should be fit to the intergalactic Hubble curve and those with (when any are found) could be included by correcting for lensing and/or by averaging in with the others to see if the weak-lensing FLRW value can be obtained. Only with enough unbiased and absorption corrected data will the FLRW Hubble curve be useful.  The intergalactic Hubble curve contains  the additional parameter $\\nu$; however, if luminous matter is essentially the whole story most matter is in galaxies and one can put $\\nu=2$ as a good approximation. Use of the intergalactic Hubble curve is then, in principle, no more involved than use of the standard FLRW Hubble curve. The division of $\\rho_0$ into  galactic and intergalactic parts for gravitational-optics purposes seems simplistic but it is certainly less simplistic than ignoring optical effects of inhomogeneities altogether as is done by using FLRW."
    },
    "9804/astro-ph9804280_arXiv.txt": {
        "abstract": "The history of the transition from a neutral intergalactic medium to one that is almost fully ionized can reveal the character of cosmological ionizing sources. In this talk I will discuss the implications for rival reionization scenarios of the rapid decline observed in the space density of quasars and star-forming galaxies at redshifts $z\\gta 3$. The hydrogen component in a highly inhomogeneous universe is completely reionized when the number of ionizing photons emitted in one recombination time equals the mean number of hydrogen atoms. At $z\\sim 5$, the local character of the UV metagalactic flux allows one to define a {\\it critical} emission rate of hydrogen-ionizing photons per unit comoving volume, ${\\dot{\\cal N}}_{\\rm ion}=10^{51.5\\pm 0.3}\\ndotunits$. Models based on photoionization by bright QSOs and/or young galaxies with star formation rates in excess of $0.3-1\\sfr$ appear to fail to provide the required number of hydrogen-ionizing photons at these redshifts by large factors. If stellar sources are responsible for keeping the universe ionized at $z\\approx 5$, the rate of star formation per unit comoving volume at this epoch must be comparable or greater than observed at $z\\approx 3$. ",
        "introduction": "The existence of a filamentary, low-density intergalactic medium (IGM) which contains the bulk of the hydrogen and helium in the universe is predicted as a product of primordial nucleosynthesis and of hierarchical models of gravitational instability with ``cold dark matter'' (CDM) (Cen \\etal 1994; Zhang \\etal 1995; Hernquist \\etal 1996). The application of the Gunn-Peterson constraint on the amount of smoothly distributed neutral material along the line of sight to distant objects requires the hydrogen component of the diffuse IGM to have been highly ionized by $z\\approx 5$ (Schneider \\etal 1991), and the helium component by $z\\approx 2.5$ (Davidsen \\etal 1996). It thus appears that substantial sources of ultraviolet photons were present at $z\\gta 5$, perhaps low-luminosity quasars or a first generation of stars in virialized dark matter halos with $T_{\\rm vir}\\gta 10^4\\,$K (Couchman \\& Rees 1986; Ostriker \\& Gnedin 1996; Haiman \\& Loeb 1997; Miralda-Escud\\`e \\& Rees 1997). Early star formation provides a possible explanation for the widespread existence of heavy elements in the IGM (Cowie \\etal 1995), while reionization by QSOs may produce a detectable signal in the radio extragalactic background at meter wavelengths (Madau \\etal 1997). Establishing the character of cosmological ionizing sources is an efficient way to constrain competing models for structure formation in the universe, and to study the collapse and cooling of small mass objects at early epochs. While the nature, spectrum, and intensity of the background UV flux which is responsible for maintaining the intergalactic gas and the Ly$\\alpha$ clouds in a highly ionized state at $z\\lta 3$ has been the subject of much debate in the last decade, it is only in the past few years that new observations have provided reliable information on the presence and physical properties of the sources and sinks (due to continuum opacities) of UV radiation in the interval $3\\lta z\\lta 5$.  In this talk I will focus on the candidate sources of photoionization at early times and on the time-dependent reionization problem, i.e. on the history of the transition from a neutral IGM to one that is almost fully ionized. The starting point of this study can be found perhaps in the simple realization that the {\\it breakthrough epoch} (when all radiation sources can see each other in the Lyman continuum) occurs much later in the universe than the {\\it overlap epoch} (when individual ionized zones become simply connected and every point in space is exposed to ionizing radiation), and that at high redshifts the ionization equilibrium is actually determined by the {\\it instantaneous} UV production rate. In the following I will adopt an Einstein-de Sitter universe ($q_0=0.5$) with $H_0= 50h_{50}\\,\\kmsmpc$. ",
        "conclusions": ""
    },
    "9804/astro-ph9804233_arXiv.txt": {
        "abstract": "\\noindent Nonlinear time-dependent calculations are being carried out in order to study the evolution of vertically-integrated models of non-selfgravitating, transonic accretion discs around black holes. In this paper we present results from a new calculation for a high-$\\alpha$ model similar to one studied previously by Honma, Matsumoto \\& Kato who found evidence for limit-cycle behaviour connected with thermal instability. Our results are in substantial agreement with theirs but, in our calculation, the disc material does not always remain completely optically thick and we include a suitable treatment for this.  We followed the evolution for several cycles and determined the period of the cycle as being about 780 seconds. Advective cooling is dominant in the region just behind the outward-moving peak of surface density.  The behaviour of this model is significantly different from what we saw earlier for low-$\\alpha$ models (which we discussed in a previous paper) and we contrast and compare the two situations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804005_arXiv.txt": {
        "abstract": "We have obtained intermediate resolution spectra of eleven candidate brown dwarf members of the Pleiades open cluster using the Keck II telescope and LRIS spectrograph.  Our primary goal was to determine the location of the ``lithium depletion edge\" in the Pleiades and hence to derive  a precise age for the cluster. All but one of our 11 program objects have radial velocities appropriate for Pleiades members, have moderately strong \\Ha\\ emission, and have spectral types M6 to M8.5 as expected from their (R-I)$_c$\\ colors.  We have constructed a color-magnitude diagram for the faint end of the Pleiades main sequence, including only stars for which high S/N spectra in the region of the lithium $\\lambda$6708$\\AA$\\ absorption line have been obtained.  These data allow us to accurately determine the Pleiades single-star lithium depletion edge at I$_{c0}$\\ = 17.80, (R-I)$_{c0}$\\ = 2.20, spectral type = M6.5.  By reference to theoretical evolutionary models, this converts fairly directly into an age for the Pleiades of $\\tau$\\ = 125 Myr.  This is significantly older than the age that is normally quoted, but does agree with some other recent estimates. ",
        "introduction": "The Pleiades was recognized in the 1980's as the best open cluster to attempt to identify brown dwarfs (Stauffer \\etal 1989; Jameson \\& Skillen 1989) because of its fortuitous combination of proximity, youth and richness.  A number of brown dwarf candidates were identified in those and subsequent papers, usually on the basis of photometry obtained from deep imaging surveys in two or more colors. A means to establish that a candidate brown dwarf was at least near, if not necessarily below, the hydrogen burning mass limit was proposed by Rebolo, Mart\\'{\\i}n and Magazz\\`u (1992; hereafter RMM).  Below about 0.065 \\MSUN, brown dwarfs should never develop core temperatures sufficient for lithium ignition, and thus for objects less massive than this we should find a lithium abundance the same as for the interstellar medium from which the object formed, independent of the object's age.  For slightly higher masses, lithium acts as an age scale because the length of time it takes for the core to reach 2.5x10$^6$\\ K is a sensitive function of mass.  Because these stars are fully convective, once the core temperature exceeds the necessary limit, the entire lithium content of the star should be exhausted rapidly and thus be reflected in an observable change in the photospheric lithium abundance.   The current generation of theoretical models make specific predictions about the time evolution of this lithium depletion boundary.  For example, D'Antona \\& Mazzitelli (1997) predict that at ages 30, 70 and 140 Myr, the lithium depletion edge should occur at 0.17 \\MSUN, 0.09 \\MSUN\\ and 0.07 \\MSUN, respectively.   Other recent models by Baraffe \\etal (1998) and Burrows \\etal (1997) make nearly identical predictions of the variation of this lithium depletion boundary with age.   Indeed, Bildsten \\etal (1997) and others have argued that the age for an open cluster derived in this manner should be better than by any other method.  The initial attempts to detect lithium in very low mass open cluster members were not successful (RMM; Marcy, Basri \\& Graham 1994; hereafter MBG94).  However, Basri, Marcy \\& Graham (1996; hereafter BMG96)  eventually detected lithium in PPL 15, a Pleiades brown dwarf candidate originally identified by Stauffer, Hamilton \\& Probst (1994).  Rebolo et al. (1996) later showed that two other Pleiades members about 1 magnitude fainter than PPL 15 also have strong lithium absorption.   BMG96 derived an age for the Pleiades based on the presence of lithium in PPL 15 but the absence of lithium in another only slightly brighter Pleiades member (HHJ3).  However, Basri \\& Mart\\'{\\i}n (1998) have recently provided evidence that suggests that PPL 15 is a nearly equal-mass spectroscopic binary and thus its individual components would be $\\sim$0.75 mag fainter.  This allows the location of the lithium boundary to be considerably fainter than had been assumed.   The published data therefore no longer constrain the age of the Pleiades nearly as well as one would like because too few stars have been measured spectroscopically in the  magnitude range of interest.  In this paper, we report on the results of a program to determine lithium abundances for a number of candidate Pleiades members with apparent magnitudes chosen to bracket the possible magnitude range within which the lithium depletion boundary might be located. ",
        "conclusions": "The goal of this project was to precisely determine the location of the lithium depletion edge in the Pleiades, and hence to determine an accurate age for the cluster.  Figure 2 shows a color-magnitude diagram for very low mass Pleiades members for which lithium data are available.  The lithium data are from Oppenheimer \\etal (1997 = Opp97), MBG94, BMG96, Rebolo \\etal (1996 = Reb96) and from this paper. The (R-I)$_c$\\ colors for the Opp97 stars are from Stauffer \\etal (1995), where we have converted these spectroscopic (V-I)$_c$\\ colors to (R-I)$_c$\\ via a relation derived from the Gliese catalog M dwarfs in Leggett (1992). The (R-I)$_c$\\ colors for some stars in Reb96 and BMG96 are from PC2 indices provided by Mart\\'{\\i}n \\etal (1996) and our calibration of PC2 vs. (R-I)$_c$.  Two points for CFHT-PL-15 are shown: (R-I)$_c$\\ = 2.24 as derived from its PC2 index, and (R-I)$_c$ = 2.41 as derived from its VO index.  Using either color, CFHT-PL-15 lies below the main sequence defined by the other stars, possibly indicating that it is a non-member.  Because it has detected lithium and a radial velocity compatible with Pleiades membership, we prefer to believe it is a member and that the inferred color is anomalous (either intrinsically or due to measurement error).  Our primary conclusions are unaffected by how we interpret this star. Finally, the dashed line in Figure 2 is a field star ZAMS derived from the M dwarf photometry provided by Leggett (1992).   The location of the stars in Figure 2 show a well-defined correlation with the measured lithium equivalent widths.  Bluer than (R-I)$_c$\\ = 2.2, none of the stars have detected lithium while redder than that color all of the measured stars have lithium.  Similarly, fainter than I$_{c0}$\\ = 17.8 all the stars have detected lithium, whereas brighter than that limit all but one of the stars have no lithium.  Finally, there is a trend for lithium equivalent width to increase going to lower inferred mass (i.e. fainter and redder).  We further believe that the dispersion in I magnitude at a given color is to a large extent real and is primarily an indication that some of the observed stars are photometric binaries. In particular, we suggest that HHJ6 (I$_{c0}$\\ = 16.93, R-I$_c$\\ = 2.18), Lick-PL1, PPL1, and CFHT-PL-12 are good candidates to be photometric binaries.  CFHT-PL-12 also has considerably stronger \\Ha\\ emission than the other Pleiades stars observed - possibly indicating it is a short-period binary or that we have observed it during a flare.   The attribution of PPL1 as a nearly equal mass binary would explain why it has strong lithium absorption despite an I$_c$\\ magnitude equal to or brighter than two other cluster members with no detected lithium.  An alternative explanation would be that there is a significant age spread in the Pleiades and these four over-luminous stars are the youngest in our sample.  By comparison to theoretical models (e.g. D'Antona \\& Mazzitelli 1997), their displacement about 0.5 mag above other stars of the same color would require them to be more than 50 Myr younger than the other stars, which we believe is unlikely (see, for example, discussions of this issue by Soderblom \\etal 1993 and Steele \\& Jameson 1995).   Based on the above interpretation of Figure 2, we determine that the single star lithium depletion edge in the Pleiades is at I$_{c0}$\\ = 17.8 $\\pm$ 0.1 or (R-I)$_{c0}$\\  = 2.20 $\\pm$ 0.05.  The uncertainty estimates are not rigorous, and arise mostly from the uncertainties in the absolute calibration of the photometry in Bouvier \\etal (1998).  For theoretical evolutionary models which incorporate realistic model atmospheres as their outer boundary condition, and hence which can predict observational colors and magnitudes for brown dwarfs, it is possible to convert directly the empirical lithium depletion edge to an age estimate for the Pleaides. In Figure 3, we plot the absolute I$_c$\\ magnitude of the lithium depletion boundary (defined here as the point where lithium has been depleted by a factor of 100) for the most recent models of Baraffe \\etal (1998) as a function of age.  To place our empirically measured point into this diagram, we assume that the cluster has (m-M)$_o$\\ = 5.60 (r $\\sim$ 130 pc), and A$_I$\\ = 0.06 (c.f. Pinsonneault \\etal 1998), leading to M(I$_c$) = 12.2 $\\pm$\\ 0.15 for the lithium depletion boundary, where we have assumed plausible but again unrigorous 1$\\sigma$\\ uncertainties of 0.1 mag for the distance modulus and 0.03 mag for the extinction and that the uncertainties add in quadrature.  The age derived in this way is then 125 $\\pm$\\ 8 Myr, where the uncertainty estimate only comes from propagating the 0.15 mag uncertainty of the boundary through the model shown in Figure 3.   In order to try to assess the model dependence of this estimate, we have also made a similar calculation for theoretical evolutionary models by Burrows \\etal (1997) and D'Antona \\& Mazzitelli (1997).   Those models do not provide R and I magnitudes, so instead we have used the (R-I)$_c$ color to estimate an I$_c$-band bolometric correction (we adopted the Monet \\etal (1992) BC$_I$ vs. (V-I)$_c$\\ relation, and a conversion from (V-I)$_c$\\ to (R-I)$_c$\\ based on data in Leggett 1992). In that manner, we estimate that the lithium depletion boundary in the Pleiades is at M(Bol) = 11.99.  Comparing this number to the predictions of the two theoretical models, we get an age for the Pleiades of 130 Myr based on the D'Antona \\& Mazzitelli (1997) calculations and 125 Myr for the Burrows \\etal (1997) models.  The most commonly quoted age for the Pleiades is of order 70-80 Myr (Patenaude 1978; Mermilliod 1981).  However, models with a relatively large amount of convective core overshoot can yield much larger ages, as was originally shown by the models of the Padova group (c.f. Mazzei \\& Pigatto 1989, who derived an age for the Pleiades of 150 Myr).  Other recent models give ages intermediate between these values (e.g. 100 Myr for Meynet, Mermilliod \\& Maeder 1993 and $\\geq$\\ 120 Myr for Ventura \\etal 1998).  Given the disagreement over the age derived from the upper main sequence turn-off, it is particularly useful to have an independent means to derive the age from low mass stars.  Using the lithium detection in PPL15 and the non-detection of lithium in HHJ3, BMG96 estimated the age of the Pleiades to be 115 Myr.  Our new result pushes this age even slightly older, but more importantly does so using many more stars and thus provides a much better defined age from the lithium data.  Finally, we note that we have chosen to use the ``traditional\" distance scale to the Pleiades (r $\\sim$\\ 130 pc).   That distance conflicts with the new Hipparcos distance to the Pleiades of about 116 pc (van Leeuwen \\& Ruiz 1997; Mermilliod \\etal 1997).  We have done this because we believe that the Hipparcos distance for the Pleiades is not correct, as has been discussed in Pinsonneault \\etal (1998) and Soderblom \\etal (1998). The zeroth order effect of simply using the Hipparcos distance and keeping everything else the same would lead to an even older Pleiades age of about 140 Myr."
    },
    "9804/astro-ph9804097_arXiv.txt": {
        "abstract": "Multiwavelength observations of high energy flare in 1996 from 3C 279 seems to favour the so called mirror model between different inverse Compton scattering models proposed as a possible explanation of gamma-ray emission in blazars. We performed kinematic analysis of the relativistic blob - mirror system and found that only part of the mirror located very close to the jet axis (very likely inside the jet cone) can re-emit soft photons which serve as a target for production of $\\gamma$-rays by relativistic electrons in the blob. Since the presence of well localized scattering mirror inside the jet is problematic, this makes problems for the mirror model. The time scale and the shape of the $\\gamma$-ray flare should reflect, in terms of the mirror model, the blob dimensions and the longitudinal distribution of relativistic electrons inside the blob. For the $\\gamma$-ray light curve of the type observed in 1996 from 3C 279, i.e. the rising time of the flare during a few days with a sharp cut-off towards the end of the flare, the density of electrons inside the blob should increase exponentially starting from the front of the blob and reach maximum towards the end of the blob. Such distribution of electrons is difficult to explain in a model of a relativistic shock moving along the jet, which would rather inject electrons more efficiently at the front of the blob with a trail of particles on its downstream side. ",
        "introduction": "About 50 blazars have been detected by the Compton Gamma Ray Observatory in the MeV - GeV  energy range (Fichtel et al. 1994, von Montigny et al. 1995, Thompson et al. 1995, Mukherjee et al. 1997), and 3 blazars, of the BL Lac type, are discovered in the TeV $\\gamma$-rays by the Whipple Observatory (Punch et al. 1992, Quinn et al. 1996, Catanese et al. 1997). These blazars can reach very high $\\gamma$-ray luminosities which are variable on time scales as short as a part of a day, in the case of optically violent variable quasars, or even several minutes, in the case of BL Lacs. These observations strongly suggest that $\\gamma$-ray emission from blazars is collimated towards the observer within a small angle as a result of relativistic motion of plasma in the jet or directional acceleration of particles.  High energy processes occurring in blazars are popularly explained in terms of the inverse Compton scattering (ICS) model in which $\\gamma$-rays are produced in ICS of soft photons by electrons in a blob moving relativistically along the jet. Different modifications of this general model mainly concern the origin of soft photons, i.e. whether they come internally from the blob in the jet (synchrotron self-Compton (SSC) model, e.g. Maraschi et al. 1992, Bloom \\& Marscher 1993), directly from the disk (e.g. Dermer et al. 1992, Bednarek et al. 1996a,b), are produced in the disk but reprocessed by the matter surrounding the disk (external comptonization (EC) model, e.g. Sikora et al. 1994, Blandford \\& Levinson 1995), or produced in the jet but reprocessed by the matter surrounding the jet (the so-called mirror model, Ghisellini \\& Madau~1996, henceforth GM). In this last paper it is mentioned that SSC model  and external comptonization of photons produced by the broad line region clouds (BLR) illuminated by the disk (EC model) may also contribute to the $\\gamma$-ray emission producing a first $\\gamma$-ray pre-flare. For the SSC model the amplitude of the $\\gamma$-ray variation is expected to be proportional to the square of the variation observed in IR-optical-UV energy range. For the EC model the $\\gamma$-ray emission should vary linearly with the low energy synchrotron emission. Such behaviour is not observed in the case of the 1996 flare from 3C 279 in which the $\\gamma$-ray variation is more than the square of the synchrotron variation. Moreover, in the $\\gamma$-ray light curve of this flare (see Fig.~1 in Wehrle et al. 1997), there is no clear evidence for a double peak structure which could eventually correspond to the first $\\gamma$-ray flare produced in terms of SSC or EC models and the second $\\gamma$-ray flare produced in terms of the mirror model. Therefore, although the SSC model can not be completely rule out, Wehrle et al. (1997) concludes that the mirror model is favourite by the multiwavelength observations of a strong flare in February 1996 from 3C 279 since it predicts $\\gamma$-ray flare with observed features.  In this paper we test the mirror model by comparing predictions of the kinematic analysis with the observational results. The possible contributions from SSC and EC models to the $\\gamma$-ray production during this flare are neglected since, as we mentioned above, there is no observational support for their importance. Simultaneous analysis of all these models will require an introduction of additional free parameters (density of electrons in the blob, the perpendicular extend of the blob, definition of the disk radiation) which are not all well constrained by the observations. ",
        "conclusions": "We discuss details of the mirror model proposed by Ghisellini \\& Madau. This model seems to be favourite by the multiwavelength observations of the $\\gamma$-ray flare in 1996 from 3C 279 (Wehrle et al. 1997). Based on the analysis of the kinematics of the emission region (a blob moving relativistically along the jet) we come to the conclusion that only relatively small part of the mirror is able to re-emit soft photons which serve as a target for production of $\\gamma$-rays. For the parameters of the $\\gamma$-ray flare observed in 1996 from 3C 279, the radius of this part of the mirror should be comparable to the longitudinal extend of the blob. It has to be of the order of $2 \\times 10^{16}$ cm in order to be consistent with the rising time of the flare. This part of the mirror should lay inside the jet cone provided that its opening angle is of the order of $\\sim 1/\\gamma$. As mentioned in Ghisellini \\& Madau (GM), the physical processes in the jet may prevent the presence of the well localized mirror inside the jet.  The calculations of density of photons re-emitted by the mirror are done by Ghisellini \\& Madau (see Fig.~2 in GM) in a time independent picture which do not take into account the dynamics of the blob. As a consequence they integrate over the parts of the mirror at distances from the jet axis  which are much larger than the maximum distance $h_{\\rm u}$ (Eq.~(\\ref{eq17})), found in our dynamical (time dependent) analysis. The photon densities seen by the blob can not be directly compared with that ones obtained by us in a time dependent version of the mirror model. Ghisellini \\& Madau results are only correct for the continuous (time independent) flow of relativistic plasma along the jet axis but overestimates the density of soft photons seen by the relativistic electrons in the blob with limited longitudinal extend. The relativistic blobs in blazars has to be confined to the part of the jet in order to produce the $\\gamma$-ray flares with the observed rising time scale.  We computed the $\\gamma$-ray light curves expected in the dynamical version of the mirror model for different distribution of relativistic electrons inside the blob and assuming that the density of electrons in the blob changes during propagation along the jet. Slowly rising $\\gamma$-ray flux with sudden cut-off towards the end of the flare, as observed in 3C 279, is obtained in the case of inhomogeneous blob with electron densities exponentially rising towards the end of the blob. Such electron distribution is difficult to understand in the popular scenario for $\\gamma$-ray production in which relativistic shock moves along the jet. It seems that such shock should rather inject relativistic electrons with high efficiencies close to the front of the blob, with the trail of electrons on its downstream side (Kirk, Rieger \\& Mastichiadis~1998). However the $\\gamma$-ray light curve expected in this case is different than observed during the flares in the blazar 3C 279.  Since $\\gamma$-rays are produced in a region which is close to the mirror, therefore the shape of the $\\gamma$-ray light curve is not very sensitive on the variations of the density of electrons during the time of propagation of the blob between the base of the jet and the mirror. Of course the absolute $\\gamma$-ray fluxes produced by the blobs with different evolutions of electron densities in time may differ significantly.  The $\\gamma$-ray light curves presented in Figs.~\\ref{fig2} show very sharp cut-offs towards the end of the flare due to our assumption on the negligible thickness of the mirror. In fact, the observed width of the peak in the $\\gamma$-ray light curve of 3C 279, of the order of $t_m\\sim 1$ day (see Fig.~1 in Wehrle et al. 1997), may be related to the time in which relativistic blob is moving though the mirror with the finite thickness. If this interpretation is correct then the thickness of the mirror has to be limited to $\\rho_{\\rm m}\\approx c t_{\\rm m} \\beta (1+\\beta)\\gamma^2 \\approx 4\\times 10^{17}$ cm which is comparable to the distance of the mirror from the base of the jet.  In this analysis we do not consider production of $\\gamma$-rays in terms of the SSC and EC models simultaneously with the mirror model since there is no clear evidence of their importance in the $\\gamma$-ray light curve and the multiwavelength spectrum observed in 1996 from 3C 279 (Wehrle et al. 1997). The $\\gamma$-ray light curves reported in Figs.~2 show only relative change of the $\\gamma$-ray flux with time. They are not straightforwardly dependent on the parameters of the blob (the magnetic field, electron density, blob perpendicular extend, disk radiation) which are not uniquely constrained by the observations. The SSC and EC models will require to fix these parameters in order to guarantee reliable comparisons."
    },
    "9804/astro-ph9804268_arXiv.txt": {
        "abstract": "Observations of galactic black hole candidates made by the instruments aboard the Compton GRO in the hard X-ray and $\\gamma$-ray bands have significantly enhanced our knowledge of the emission properties of these objects. Understanding these observations presents a formidable challenge to theoretical models of the accretion flow onto the compact object and of the physical mechanisms that generate high-energy radiation.  Here we summarize the current state of observations and theoretical interpretation of the emission from black hole candidates above 20 keV.  The all-sky monitoring capability of BATSE allows, for the first time, nearly continuous studies of the high-energy emission from more than a dozen black hole candidates.  These long-term datasets are particularly well-suited to multiwavelength comparison studies, from the radio upward in frequency (Zhang et al. 1997a, these proceedings). Energy spectral evolution and/or spectral state transitions have been observed from many of the black hole candidates. Moderately deep searches of the galactic plane suggest a deficit of weak $\\gamma$-ray transients.  Such population studies have implications for the origin of black hole binaries and the nature of accretion events.  Observations above 50 keV from OSSE demonstrate that in the $\\gamma$-ray band there exist two spectral states that appear to be the extensions of the X-ray low (hard) and high (soft), or perhaps very high, states. The former state, the ``breaking'' state, cuts off with e-folding energy $\\sim$100 keV and has its peak luminosity near this energy; thus substantial corrections need to be made to historical estimates of the bolometric luminosity of black holes in the ``low'' state.  In contrast, in the X-ray high (soft) state, the luminosity peaks in the soft X-rays and the spectrum extends with an unbroken power law, even up to energies above 500 keV in some cases. COMPTEL has detected emission above 750 keV from Cyg X-1 and the transient GRO~J0422+32.  In both cases the data suggest that an additional weak, hard spectral component is required beyond that observed by OSSE at lower energies, although the precise spectral form is yet to be determined.  The breaking $\\gamma$-ray spectrum can be well modeled by Comptonization of soft photons from the accretion disk in a hot thermal plasma. However, recent studies of the combined X-ray and $\\gamma$-ray spectrum of Cyg~X-1 and GX339--4 cast severe doubts on the simple geometry of a hot corona overlying a thermal accretion disk. Furthermore, timing studies of the former source are inconsistent with spectral formation by Compton scattering in a uniform, compact hot cloud, suggesting instead a decline in electron density with increasing radius.  The power-law $\\gamma$-ray spectral state creates more significant theoretical challenges, particularly in explaining the lack of a break at energies exceeding the electron rest mass. It has been suggested that in the X-ray high (soft) state, the high-energy emission arises from bulk-motion Comptonization in the convergent accretion flow from the inner edge of the accretion disk.  Such a process can conceivably generate the $\\gamma$ ray spectrum extending without a cutoff, if the accretion rate approaches that of Eddington. ",
        "introduction": "The most reliable evidence for the presence of a black hole in a binary system comes from determination of a mass function through optical measurements of the radial velocity of the companion star.  If the resulting lower limit on the mass of the compact object exceeds 3$\\Msun$, the upper limit for the mass of a stable neutron star based on current theory, then one can reasonably assume that the compact object is a black hole.  There are at least nine X-ray binary systems with minimum mass estimates exceeding 3$\\Msun$, of which three (Cyg~X-1, GRO~J0422+32, and GRO~J1655--40) have been clearly detected by GRO instruments.  Other objects are identified as BHCs based on the similarity of their high-energy spectra and rapid time variability to those of Cyg~X-1.  Such classification is, of course, somewhat tenuous.  Before neutron stars and black holes can be reliably distinguished based on their X-ray and $\\gamma$-ray spectra, the full range of spectral forms from both classes must be observed and characterized. Extensive knowledge of the X-ray emission of these objects has accumulated in the literature, but the broad nature of the $\\gamma$-ray emission is only now coming to light, with the high sensitivity of current-generation instruments.  The instruments of the Compton GRO have made extensive observations in the hard X-ray and $\\gamma$-ray bands of galactic black hole candidates (BHCs).  With its all-sky capability, BATSE has monitored emission on a nearly continuous basis from at least three persistent sources (Cyg~X-1, 1E1740.7--2942, GRS~1758--258) and eight transients (GRO~J0422+32, GX339--4, N Mus 1991, GRS~1716--249, GRS~1009--45, 4U~1543--47, GRO~J1655--40, and GRS~1915+105).  Lightcurves are presented below in Fig. \\ref{lightcurve}.  OSSE has made higher-sensitivity, pointed observations of all of these sources, spectra of which appear below in Fig. \\ref{two_state}. COMPTEL has detected emission above 750 keV from Cyg~X-1 and GRO~J0422+32.  To date, there have been no reported detections of galactic BHCs by EGRET.  The French coded-aperture telescope Sigma on the Russian Granat spacecraft has imaged at least a dozen BHCs, including most of those in the list above, but with the addition of TrA~X-1, GRS~1730--312, and GRS~1739--278.  The latter two objects were weak transients discovered during a multi-year survey of the galactic center region and have been classified as BHCs by their outburst lightcurves and the hardness of their spectra (Vargas et al. 1997).  In this survey, Sigma regularly detected the persistent, variable sources 1E1740.7--2942 and GRS~1758--258, both of which are classified as BHCs on spectral grounds.  The most striking result from Sigma observations of BHCs is the detection of broad spectral features below 500 keV from 1E1740.7--2942 (Sunyaev et al. 1991, Bouchet et al. 1991, Churazov et al. 1993, Cordier et al. 1993) and N Mus 1991 (Goldwurm et al. 1992, Sunyaev et al. 1992).  These features have been interpreted as thermally broadened and red-shifted annihilation radiation from the vicinity of the compact object. ",
        "conclusions": ""
    },
    "9804/astro-ph9804118_arXiv.txt": {
        "abstract": "We present results of \\ASCA\\ deep exposure observations of the hardest X-ray source discovered in the \\ASCA\\ Large Sky Survey (LSS) project, designated as AX~J131501$+$3141. We extract its accurate X-ray spectrum, taking account of the contamination from a nearby soft source (AX~J131502$+$3142), separated only by 1$'$. AX~J131501$+$3141 exhibits a large absorption of $N_{\\rm H} = (6^{+4}_{-2})\\times 10^{22}$ \\NHUNIT\\ with a photon index $\\Gamma = 1.5^{+0.7}_{-0.6}$.  The 2--10 keV flux was about $5\\times 10^{-13}$ \\FLUXUNIT\\ and was time variable by a factor of 30\\% in 0.5 year.  From the highly absorbed X-ray spectrum and the time variability, as well as the results of the optical follow-up observations (\\cite{Akiyama98}), we conclude that AX~J131501$+$3141 is a type 2 Seyfert galaxy.  Discovery of such a low flux and highly absorbed X-ray source could have a significant impact on the origin of the cosmic X-ray background. ",
        "introduction": "Since the discovery of the Cosmic X-Ray Background (CXB) more than 30 years ago (\\cite{Giacconi62}), its origin has been a long standing puzzle.  With the \\ROSAT, $\\sim$70\\% of the CXB below 2 keV has been resolved into point sources, more than 60\\% of which are type 1 active galactic nuclei (AGNs) (\\cite{Vikhlinin95}; \\cite{Hasinger96}; \\cite{McHardy98}; \\cite{Hasinger98}).  Origin of the CXB above the 2 keV band, however, is less clear due to the absence of the imaging instrument in this energy band.  One problem in the hard X-ray band, often referred as the spectral paradox (e.g. \\cite{Fabian92}), is that the X-ray spectrum of the CXB in the 2--10 keV band is harder than that of the typical type 1 AGN, which is presumably the main contributor to the CXB below 2 keV.  The 2--10 keV X-ray spectra of type 1 Seyfert galaxies (most of the bright AGNs) are approximated by a power-law with a mean photon index of 1.7 (\\cite{Mushotzky93}), which is significantly steeper than that of the 2--10 keV CXB of about 1.4 (\\cite{Gendreau95}).  This fact implies that the origin of the CXB above 2 keV differs, at least in part, from that below 2 keV.  In addition, 20\\% of the total energy of the CXB is contained in the 2--10 keV band, whereas only a few percent of the CXB is contained below 2 keV (see review by Fabian \\& Barcons (1992), Hasinger (1996)).  Thus, the 2--10 keV band would be an essential energy range to solve the origin of the CXB.  \\ASCA\\ is the first satellite with the capability of the hard X-ray (up to 10keV) imaging and spectroscopy, hence is presently the best satellite to investigate the CXB in the 2--10 keV band.  In the \\ASCA\\ Large Sky Survey project (LSS: \\cite{Inoue96}; \\cite{Ueda98a}), a continuous field of 7 deg$^2$ near the north galactic pole was surveyed with a sensitivity higher than any previous surveys in this energy band.  Ueda \\etal\\ (1998a) resolved a significant fraction of the CXB, about 30\\% of the CXB, into discrete sources at a sensitivity limit of $F_{\\rm X}\\sim 10^{-13}$ \\FLUXUNIT\\ (2--10 keV).  The mean photon index in the 2--10 keV band for these resolved sources ($F_{\\rm X}=$(1--4)$\\times 10^{-13}$ \\FLUXUNIT\\ in 2--10 keV) was found to be $\\Gamma = 1.5\\pm 0.2$.  This result is consistent with the idea that the photon index approaches to that of the CXB, $\\Gamma\\sim 1.4$, as the source flux decreases. However, due to limited photon statistics, the spectral information of the resolved sources was too poor to address the nature of individual X-ray sources.  The hardest source in the LSS (hereafter we refer it as the ``LSS hardest source'') was found to show a photon index of $\\Gamma\\sim -0.2$ with no correction of an absorption (\\cite{Ueda96}). However, it is unclear whether the apparent hard spectrum is due to a large absorption or due to flatness of the intrinsic spectrum.  The LSS hardest source, which was found in an unbiased survey, would provide us a good opportunity to investigate the nature of faint and hard sources which could significantly contribute to the CXB above 2 keV. Hence, we have performed follow-up \\ASCA\\ and optical observations on the LSS hardest source. This paper reports results of the \\ASCA\\ deep exposure observations, while those of the optical observations are given by Akiyama \\etal\\ (1998). ",
        "conclusions": "We extracted the accurate spectrum of the LSS hardest source, AX~J131501$+$3141, taking account of the contamination from the nearby soft source, AX~J131502$+$3142, from which no significant X-ray emission was found in the LSS (Ueda 1996). We found that AX~J131501$+$3141 exhibits a large absorption of $N_H = (6^{+4}_{-2})\\times 10^{22}$ \\NHUNIT\\ with a photon index $\\Gamma = 1.5^{+0.7}_{-0.6}$. It showed a long-term time variability between two observations separated by 0.5 year. While the photon index of AX~J131501$+$3141 is consistent with the canonical value of type 1 AGNs (e.g., Mushotzky 1993), its absorption column density is larger than that of typical type 1 AGN by more than an order of magnitude, although a part of type 1 AGNs, about 10\\% of them (\\cite{Schartel97}), shows a column density larger than $5\\times 10^{22}$ \\NHUNIT.  It is important that this source is selected fully unbiasedly.  Hence, its X-ray properties should provide a key to understand the general nature of the missing hard X-ray populations which constitute the CXB above 2 keV. Two major possibilities have been proposed to account for the apparent hard spectrum of the CXB: one is to introduce large absorptions of sources (e.g., \\cite{Awaki91}), and the other is to consider populations of sources with intrinsically flat spectra (e.g., \\cite{Morisawa90}; \\cite{Matteo97}). Our results of the LSS hardest source strongly suggest that highly absorbed sources play an important role in considering the origin of the hard X-ray background.  The large absorption of $6\\times 10^{22}$ \\NHUNIT, the photon index of $\\Gamma\\sim 1.5$, and the time variability are common properties seen in type 2 Seyfert galaxies.  In fact, systematic studies of type 2 Seyfert galaxies by Awaki \\etal\\ (1991) and Ueno (1996) revealed that they commonly show large absorptions of $\\sim 10^{23}$ \\NHUNIT\\ and photon indices of 1.5--1.7. Akiyama \\etal\\ (1998) found one bright optical galaxy with $B=17.25$ mag near the center of the X-ray error circle of 0.5$'$ radius in the optical follow-up observations. No other optical source with the flux larger than $B=22.4$ mag is found in the error circle. Akiyama \\etal\\ (1998) performed spectroscopic observations of the bright galaxy and found that ratios of emission lines are similar to those found in type 2 Seyfert galaxies. The redshift of this galaxy was determined to be 0.07. From the redshift, the observed flux in the 2--10 keV band, $5\\times 10^{-13}$ \\FLUXUNIT, can be converted to the absorption corrected luminosity of $L_{\\rm X}\\sim 2\\times 10^{43}\\ {\\rm erg~s}^{-1}$.  This luminosity is consistent with those of Seyfert galaxies. Thus, we identify AX~J131501$+$3141 found in the unbiased X-ray survey as a type 2 Seyfert galaxy. Using the LogN-LogS relation in Hasinger \\etal\\ (1998), we estimate that the chance coincidence between AX J131501$+$3141 and AX J131502$+$3142 is $\\sim$ 3\\%. However, these two sources have probably no physical correlation, because AX J131502$+$3142 is likely to be a QSO\\footnote{ In the optical imaging observations of R- and B-band by Akiyama \\etal\\ (1998), we found two point-like optical sources located at about 10$''$ north from the center of the error circle with 30$''$ radius for the soft source (AX J131502$+$3142). They are very close to each other and their magnitudes are comparable. Total magnitudes of the two sources are B=20.8 mag and R=19.6 mag; B-R color is 1.21. Since both the optical color and the optical to soft X-ray flux is consistent with that for type-1 AGNs, it is possible that one of them is a quasar which is also responsible for AX J131502$+$3142. However there are fainter optical sources (R$\\gtrsim$22) in the error circle and the optical identification for AX J131502$+$3142 is not clear yet. } which is more distant than the new type 2 Seyfert AX J131501$+$3141.  Awaki (1991), Madau, Ghisellini \\& Fabian (1994) and Comastri \\etal\\ (1995) predicted that the combination of type 1 and type 2 AGNs can reproduce the CXB spectrum, based on the unified AGN scheme (e.g., \\cite{Antonucci93}).  In the scheme, type 1 and type 2 AGNs are essentially the same objects, observed from different viewing angle. These type 2 AGNs, which exhibit apparently fainter and harder X-ray spectra than those of type 1 AGNs, should become detectable as the detector sensitivity increases. Although we have examined only one sample from the LSS at this moment, the result is encouraging not only for the unified AGN scheme, but also for solving the origin of the CXB."
    },
    "9804/astro-ph9804162_arXiv.txt": {
        "abstract": "We present observations of NGC 4038/39 in the [\\ion{C}{2}] 158 \\micron \\/ fine structure line taken with the MPE/UCB Far-infrared Imaging Fabry-Perot Interferometer (FIFI) on the KAO. A fully sampled map of the galaxy pair (without the tidal tails) at 55\\arcsec \\ resolution has been obtained. The [\\ion{C}{2}] emission line is detected from the entire galaxy pair and peaks at the interaction zone. The total [\\ion{C}{2}] luminosity of the Antennae is $L_{\\rm [C II]} = 3.7 \\times 10^{8} L_{\\sun}$, which is about 1\\% of the far-infrared luminosity observed with IRAS. The main part of the [\\ion{C}{2}] emission probably arises from photodissociation regions (PDRs), and a minor fraction may be emitted from \\ion{H}{2} regions. A small part of the [\\ion{C}{2}] emission comes from standard cold neutral medium (CNM); however, for high temperature ($T \\sim 100$~K) and high density ($n_{\\rm H} \\sim  200$~cm$^{-3}$) about one third of the observed [\\ion{C}{2}] emission may originate from CNM. From PDR models we derive densities of the order of $\\sim 10^{5}$~cm$^{-3}$ and far-UV (FUV) intensities of $460\\chi_{\\circ}$, $500\\chi_{\\circ}$, and $240\\chi_{\\circ}$ for the PDRs in the interaction zone, NGC~4038, and NGC~4039, respectively. However, PDRs with densities of the order of $\\sim 10^{2}$~cm$^{-3}$ and FUV intensities of the order of $\\sim 100\\chi_{\\circ}$ could also explain the observed [\\ion{C}{2}] emission. The minimum masses in the [\\ion{C}{2}] emitting regions in the interaction zone and the nuclei are a few $\\times 10^{7}~M_{\\odot}$. A comparison with single dish CO observations of the Antennae shows a [\\ion{C}{2}] to CO intensity ratio at the interaction zone a factor of 2.6 lower than usually observed in starburst galaxies, but still a factor of about 1.3 to 1.4 higher than that at the nuclei of NGC 4038/39. Therefore, no global starburst is taking place in the Antennae. [\\ion{C}{2}] emission arising partly from confined starburst regions and partly from surrounding quiescent clouds could explain the observed [\\ion{C}{2}] radiation at the interaction zone and the nuclei, though the star formation activity toward the nuclei is lower. Accordingly there are small confined regions with high star formation activity in the interaction zone and with a lower star formation activity in the nuclei. This supports the high density and high FUV intensity for the PDRs in the interaction zone and the nuclei. ",
        "introduction": "The galaxy pair NGC 4038/39 (Arp 244) is an interacting system in an early stage of merging at a distance of about 21 Mpc from our own galaxy. On long-exposed images in the optical (e.g. Arp 1966) the interaction is clearly visible because of the tails (``Antennae'') emerging from two uniformly luminous, partly overlapping ovals and because of the dwarf galaxy that appears to have formed at the tip of the southern tail through the interaction (Zwicky 1956, Schweizer 1978, Mirabel, Dottori, \\& Lutz 1992). The tails contain about 70~\\% of the total amount of \\ion{H}{1} in the system (van der Hulst 1979). Short-exposed images (Laustsen, Madsen, \\& West 1987) reveal hints of the interaction on somewhat smaller scales; the distorted arrangement of H$\\alpha$ knots and the velocity distribution of the individual knots lead Rubin \\etal \\ (1970) to conclude there is an interaction of two rotating galaxies. Computer simulations carried out in the classical paper of Toomre \\& Toomre (1972) and later by Barnes (1988) can account for the present morphological appearance of the system quite well by assuming an interaction of two rotating spiral galaxies.  Spectra of the nuclei of both galaxies taken in the optical range do not look like pure starburst spectra but consist of a composition of early-type stars and late giants (Keel \\etal \\ 1985). The detection of bright near-infrared peaks at the nuclei lead Bushouse \\& Werner (1990) to the same result. Their images of the Antennae in J- and R-band and in H$\\alpha$ show also the same pattern of bright knots in the surroundings of NGC 4038 and in the bridge connecting both galaxies. The Antennae system as a whole shows a relatively low star forming efficiency according to the ratio $L_{\\rm IR} /M({\\rm H}_2) \\approx 8.55$ measured by Young \\etal \\ (1986); it is only a factor of 3 higher than in the Milky Way. The ratio determined by Sanders \\& Mirabel (1985) is almost twice as high. However, they observed a smaller region in CO and therefore probably underestimated the molecular mass. Comparison of the Antennae with the sample of interacting and isolated galaxies of Young \\etal \\ (1986) shows NGC 4038/39 to have characteristics more like isolated galaxies.  Measurements in the radio continuum at 1.5~GHz and 4.9~GHz of NGC~4038/39 (Hummel \\& van der Hulst 1986) reveal a number of discrete knots which coincide in general with H$\\alpha$ knots, and an underlying diffuse component. This diffuse component has a steep spectral index on average which indicates non-thermal emission, and the peak of the diffuse radio emission is at the dust patch near the overlapping region. The discrete radio knots account for roughly 35~\\% of the total radio emission and have a spectral index of $\\alpha \\approx -0.5$ on average, probably due to a thermal contribution (Hummel \\& van der Hulst 1986).  Interferometric observations of the Antennae in CO ($1 \\to 0$) by Stanford \\etal \\ (1990) show three main concentrations of CO emission. Two are associated with the nuclei and the third with the interaction zone. The overlap region is the strongest CO source and contains $\\approx 10^{9}~M_{\\sun}$ of gas, roughly as much H$_{2}$ as both nuclei together. Based on 10$\\mu$m and H$\\alpha$ data, the authors have calculated a star forming rate of $5 M_{\\sun}$~yr$^{-1}$, and consequently the life time of the molecular gas is $2 \\times 10^{8}$~yr. Single dish observations in CO were made at the interaction zone and both nuclei of the Antennae by Aalto \\etal \\ (1995). The ratio of the emission lines of $^{12}$CO and $^{13}$CO measured in the nuclei and the overlapping region of the Antennae is similar to that found in the central regions of ``normal'' starburst galaxies (Aalto \\etal \\ 1995).  An excellent tracer of star formation activity in galaxies is the strong [\\ion{C}{2}] 158~$\\mu$m $^{2}$P$_{3/2}\\to \\, ^{2}$P$_{1/2}$ fine structure line which arises mainly from photodissociation regions (PDRs) created by far-ultraviolet photons from hot young stars impinging on nearby dense interstellar clouds (Crawford \\etal \\ 1985, Stacey \\etal \\ 1991). In combination with CO and FIR observations, the [\\ion{C}{2}] emission can be used with PDR models (Tielens \\& Hollenbach 1985, Wolfire, Hollenbach, \\& Tielens 1989, Wolfire, Tielens, \\& Hollenbach 1990) to derive densities and far-UV intensities and estimates of the star formation activity. Extragalactic surveys of [\\ion{C}{2}] emission (Crawford \\etal \\ 1985, Stacey \\etal 1991) concentrated mainly on nuclei while more recent observations have imaged individual galaxies to study the distribution of [\\ion{C}{2}] emission on large scales. [\\ion{C}{2}] images of M83 (Geis \\etal \\ 1998) and NGC~6946 (Madden \\etal \\ 1993) demonstrate that the emission is extended at least over the full optical extent and often follows the distribution of the FIR and CO within the disk of the galaxies. In the case of NGC~6946 [\\ion{C}{2}] emission beyond the optical extent of the galaxy has been found. This [\\ion{C}{2}] emission has been attributed to diffuse gas and not to PDRs. Therefore we also estimate the possible contribution of [\\ion{C}{2}] emission from neutral atomic gas and from ionized gas in NGC~4038/39. Because of its relative proximity, NGC~4038/39 is a unique source for carrying out spatially resolved measurements of the [\\ion{C}{2}] line in an interacting system. We present the results of our imaging spectroscopy study of the [\\ion{C}{2}] line in NGC~4038/39 and compare them with observations obtained with ISO. ",
        "conclusions": "We present a map of NGC 4038/39 in the [\\ion{C}{2}] 158 \\micron \\ fine structure line. [\\ion{C}{2}] emission is detected over the optical extent of the system of galaxies and peaks at the interaction zone. The total luminosity of the [\\ion{C}{2}] line is $3.7 \\times 10^{8} L_{\\sun}$ which is about 1\\% of the FIR luminosity of the Antennae. Only a negligible fraction of the observed [\\ion{C}{2}] emission can originate in the WNM, if conditions are similar to Galactic atomic clouds. Under normal conditions the [\\ion{C}{2}] emission from standard CNM makes only a small contribution to the total [\\ion{C}{2}] emission, however it may rise to $\\slantfrac{1}{3}$ of the total [\\ion{C}{2}] emission for individual positions. Only minor fractions of the [\\ion{C}{2}] emission at the interaction zone and the nuclei can arise from \\ion{H}{2} regions. PDRs are the dominant source for the [\\ion{C}{2}] emission. We estimate minimum hydrogen masses associated with the [\\ion{C}{2}] emitting region of $1.9 \\times 10^{8} M_{\\sun}$ for the entire merging system and $6.8 \\times 10^{7} M_{\\sun}$, $3.2 \\times 10^{7} M_{\\sun}$, and $3.7 \\times 10^{7} M_{\\sun}$ within one beam centered at the interaction zone, NGC~4038, and NGC~4039, respectively. Assuming a single emission component in the beam we derive a density of the [\\ion{C}{2}] emitting gas in PDRs of $1 \\times 10^{5}$ cm$^{-3}$ for the interaction zone and of $2 \\times 10^{5}$ cm$^{-3}$ and $1 \\times 10^{5}$ cm$^{-3}$ for NGC~4038 and NGC 4039, respectively, and a far-UV intensity of $450 \\chi_{\\circ}$ for the interaction zone, $500 \\chi_{\\circ}$ for NGC~4038, and $250 \\chi_{\\circ}$ for NGC~4039. The derived beam filling factor of the emission from the PDRs is 20\\% for the interaction zone and 10--15\\% for the nuclei. However, the PDR model also allows a solution with low density ($\\sim 10^{2}$~cm$^{-3}$), high beam filling factor ($\\sim 50$\\%), and low FUV intensity ($\\sim 100 \\chi_{\\circ}$). The low, beam averaged [\\ion{C}{2}]/CO ratio of 2350 toward the interaction zone and the even lower ratios at the nuclei indicate that no global starburst is going on either in the area surrounding the interaction zone or in the nuclear region. The high-excitation lines observed with ISO SWS which trace the starburst must therefore arise from a small, confined region in the interaction zone. This result is also supported from the observations with ISOCAM. Therefore on the scale of our [\\ion{C}{2}] beam a single emission component for the PDR is only a crude approximation. Using interferometric and single dish CO observations and an expected [\\ion{C}{2}]/CO ratio for starburst regions and for quiescent clouds, we constructed a two-component model consisting of a confined starburst region and of molecular clouds enveloping the starburst. From this model we find that the [\\ion{C}{2}] emission at the interaction zone originates partly from confined starburst regions and partly from surrounding quiescent clouds. If we apply this model to the nuclei we get also enhanced star formation activity in NGC~4039 with a low [\\ion{C}{2}]/CO ratio for the quiescent clouds but only moderate star forming activity in NGC~4038. This two-component model supports the high-density solution for the PDRs.  Future investigation of the Antennae in the [\\ion{N}{2}] 205 \\micron \\ fine structure line would be very helpful to further disentangle the origins of the [\\ion{C}{2}] line. Also observations at higher spatial resolution in the FIR regime (e.g. with FIRST and SOFIA) would be a great step forward to investigate this and other spatially very complex objects in more detail."
    },
    "9804/astro-ph9804024_arXiv.txt": {
        "abstract": "Understanding the properties of interstellar turbulence is a great intellectual challenge and the urge to solve this problem is partially motivated by a necessity to explain the star formation mystery. This review deals with a recently suggested inversion technique as applied to atomic hydrogen. This technique allows to determine 3D turbulence statistics through the variations of 21~cm intensity. We claim that a radio interferometer is an ideal tool for such a study as its visibility function is directly related to the statistics of galactic HI. Next, we show how galactic rotation curve can be used to study the turbulence slice by slice and  relate the statistics given in galactic coordinates and in the velocity space. The application of the technique to HI data reveals a shallow spectrum of the underlying HI density that is not compatible with a naive Kolmogorov picture. We show that the random density corresponding to the found spectrum tends to form low contrast filaments that are elongated towards the observer. ",
        "introduction": "The properties of the interstellar medium strongly suggest that it is turbulent. Here turbulence is understood  as unpredictable temporal behavior of nonlinear systems as preached by J. Scalo (1985, 1987).  The importance of turbulence in molecular clouds and its relation to star formation has long been appreciated (Dickman 1985). Recent progress in numerical simulations of molecular cloud dynamics (see Ostriker, this volume) indicates the intrinsic connection between the turbulence in different phases of the interstellar medium (McKee \\&  Ostriker 1977). In what follows we shall mostly discuss the turbulence in atomic hydrogen (HI), although the formalism presented here is applicable to other spectral lines.   Statistical description is a nearly indispensable strategy when dealing with turbulence and a big advantage of statistical techniques is that they extract underlying regularities of the flow and reject incidental details. Kolmogorov notion of energy cascade from large to small scales has been proved an extremely valuable concept and Kolmogorov spectrum of turbulence has been measured in various media.    At the same time, astrophysical turbulence, unlike that in incompressible fluids, is a much more complicated phenomenon, and therefore one cannot {\\it a priori} hope that Kolmogorov's (1941) description is adequate (cf. Armstrong  et al.  1995). Energy injection at small scales, shocks, compressibility may make interstellar turbulence spectrum much more informative, and we should expect to see deviations from the boring $-11/3$ slope.  The advantage of using 21~cm emission data  is that a continuum of separations between data points is available. This property is shared by  diffuse emission in other spectral lines, but 21~cm measurements allow to disregard dust adsorption. As our review deals with HI studies within the galactic plane we do not discuss in detail interesting results obtained for HI in Large Magelanic Clouds (Spicker \\& Feitzinger 1988a,b). To avoid possible misunderstanding we should stress that Spicker \\& Feitzinger (1988a,b) deal with velocity fluctuations, while only intensity fluctuations are available when one studies HI turbulence in galactic plane. Statistics of random velocity and density fields may be different and therefore a direct comparison of the results obtained for these fields may be misleading.  Being limited in space we refer the interested reader to the earlier reviews on interstellar turbulence, among which the one by Dickman (1985) can serve as an excellent introduction to the basic statistical techniques. A more advanced reader will enjoy a thoughtful analysis of problems associated with the statistical analysis of observational data given in Houlahan \\& Scalo (1990). Important aspects of the statistical analysis are discussed, for instance, by Dickman \\& Kleiner (1985), Roy \\& Joncas (1985) Perault  et al.  (1986), O'Dell \\& Casta\\~{n}eda (1987),  Rickett (1988), Van Langevelde et al. (1992)  Kitamura  et al.  (1993), Meisch \\& Bally (1994), Armstrong, Rickett, \\& Spangler (1995), Wallin, Watson, \\& Wyld (1998) and by the contributors to the present volume. A brief discussion of the very early studies of interstellar statistics can be found in Lazarian (1992).  Studies of interstellar turbulence frequently deal with samples which are not statistically homogeneous (see Miesch \\& Bally 1994). Indeed, whenever individual molecular complexes are studied, the statistics (especially at large separations) may be dominated by regular gradients rather than the random component. To eliminate the inhomogeneous component, various types of spatial filtering are used (see  Zurfleh 1967). These problems are alleviated for HI studies, since molecular clouds tend to be localized objects in sharp contrast to more pervasive distribution of atomic hydrogen.  Further on we shall deal with the two point structure functions and power spectra. Naturally, one cannot place pickup devices at different points of HI. Instead, only the 21-cm intensity fluctuations with pointwise emissivity integrated along the lines of sight are available. It is obvious, that given a statistical description of the transparent emitting astrophysical medium, it is possible to predict statistical properties of the observable diffuse emission (Kaplan \\& Pikelner 1970), which would correspond to the solution of the {\\it forward problem}. However, more important is to solve the {\\it inverse problem}, namely, to deduce the 3D statistics of HI from observations. These issues are dealt with in sections 2 and 3. In section 4 we discuss the application of the technique to interferometric data.   The galactic rotation curve allows one to study turbulence slice by slice. This slicing, however, is far from trivial (see section 4). Indeed, topologically disconnected blobs of HI can overlap in the velocity space if their velocities are the same. We show that interferometric study can potentially provide the information about both random density and velocity fields.  Addressing the issue of HI topology we show that HI with the measured spectrum of density fluctuations forms low contrast filaments (section 5) and these filaments are elongated towards the observer due to the presence of velocity fluctuations. ",
        "conclusions": "The following are the principal conclusions of this review:  \\begin{itemize} \\item The application of a newly developed inversion technique reveals a shallow spectrum of HI density in galactic plane. The spectrum is different from the Kolmogorov one, which may be indicative of non-trivial physics involved.  \\item The spectrum of HI density is anisotropic in velocity space with velocity fluctuations altering the statistics along the line of sight. The degree of anisotropy and the spectrum in velocity space depend on the scale under study. The spectrum become isotropic for scales larger than 200-300~pc.  \\item Interferometers are useful tools for studying HI turbulence in galactic plane and may provide the spectrum of 3D random density field and useful information on random velocity field.  \\item Atomic hydrogen with Gaussian density and the shallow spectrum corresponding to observations forms low contrast filaments. These filaments are anisotropic in the velocity space and directed towards an observer. \\end{itemize}"
    },
    "9804/astro-ph9804206_arXiv.txt": {
        "abstract": "As part of a continuing study of the effect of cluster environment on the star formation properties of galaxies, we have undertaken an H$\\alpha$ objective prism survey of the nearby cluster, Abell 1060. We detect 33 galaxies in emission, 24 of which are cluster members. We present new radial velocity measurements and  H$\\alpha$+[\\ion{N}{ii}] equivalent widths and fluxes for a number of these galaxies.  We distinguish between galaxies with diffuse and compact emission, the latter having been associated in previous work with a disturbed morphology of the galaxy and most likely resulting from tidally-induced star formation from galaxy--galaxy or cluster--galaxy interactions.  The fraction of cluster spirals in Abell 1060 detected with compact emission agrees with the expected fraction for a cluster of its richness, as derived from results of a previous survey of 8 clusters. Some of the detected cluster early-type spirals exhibit anomalously high global H$\\alpha$ equivalent widths, as compared to galaxies of similar type in the field. ",
        "introduction": "\\label{resint} \\noindent  The effect of cluster environment on the star formation properties of galaxies has long been a matter of debate. While some studies have suggested a lower star formation rate for cluster spirals as compared to the field (e.g. \\citeNP{gisler,dress85}), other work has suggested a similar or enhanced rate, particularly for early-type spirals(e.g. \\citeNP{kenn84,gav91,moss93,enacs}). With the discovery that in distant rich clusters there is a high fraction of blue, star-forming galaxies, often with unusual morphology suggestive of mergers and tidal interactions (e.g. \\citeNP{lavhen86,tom88,couch94}), there is renewed interest in tidally-induced star formation by mergers and interactions in nearby clusters. In order to address these issues, Moss, Whittle and co-authors have completed an objective-prism survey of eight nearby clusters of galaxies to detect global \\mbox{H$\\alpha$+[\\ion{N}{ii}]} emission as an indicator of the current rate of massive star formation. The survey technique is described by \\citeN{mwi}, hereafter MWI, and initial results have been discussed by \\citeN{moss93} (see also \\citeNP{moss90,moss95,moss96,moss97}). We have extended this survey to a ninth cluster, the Hydra I cluster, Abell 1060.  Abell 1060 is one of the nearest of the Abell clusters, at a redshift of $z\\sim$0.01, and is the nearest large cluster beyond the Virgo and Fornax clusters. It has a high spiral fraction (e.g. \\citeNP{solanes92}) and is a relatively poor cluster, with a low density intracluster medium \\cite{lowmush} and low X-ray luminosity \\cite{edge91}. Since it is the nearest of the clusters surveyed by us so far, it can be surveyed to a fainter limit in absolute magnitude. However, its proximity means that it has a large projected diameter on the sky, with one Abell radius, \\mbox{$1\\, r_{A}=2\\fdg3=1.5\\, h^{-1}$ \\rm{Mpc}} (where $h$ is defined in terms of the Hubble constant $H_{0}=100h\\ \\rm{km\\ s}^{-1}\\rm{Mpc}^{-1}$). Whereas other clusters were surveyed in a region of radius 1.5 $r_{A}$, the photographic plate size restricted survey of Abell 1060 to a region of radius somewhat less than one Abell radius (see \\S \\ref{platty}). \\citeN{r89}, hereafter R89, presents a catalogue of 581 galaxies in the cluster area, which contains a sample which is complete to the magnitude limit $V_{25}=16.65$, within 2\\degr \\/ of the cluster centre. This is a convenient complete sample for the present H$\\alpha$ survey, extending to a fainter apparent magnitude than the Zwicky Catalogue used to define samples for other clusters.  Cluster properties are summarised in Table \\ref{Hydratab}. The (B1950.0) position of the central cluster galaxy NGC 3311 is given as the cluster centre in columns 2 and 3. The mean heliocentric radial velocity $\\langle v_{\\odot}\\rangle$ and velocity dispersion $\\sigma$ determined using $N_{gal}$ individual galaxy redshifts are given in columns 4, 5, and 6 \\cite{bird94}. The Abell richness class \\cite{aco} is given in column 7, the Bautz-Morgan and Rood-Shastry type classes are given in columns 8 and 9 respectively \\cite{bm,strood}.   \\begin{table*}  \\centering  \\caption{\\label{Hydratab} Cluster properties} \\begin{tabular}{ll@{\\hspace{1.0ex}}l@{\\hspace{1.0ex}}ll@{\\hspace{1.0ex}}l@{\\hspace{1.0ex}}lcccccc}  \\\\ \\hline \\\\ Name & \\multicolumn{6}{c}{R.A.~~(1950.0)~~Dec.} & $\\langle v_{\\odot}\\rangle$ & \\(\\sigma\\) & $N_{gal}$ &Richness&\\multicolumn{2}{c}{Type Class} \\\\ \\cline{12-13} &\\multicolumn{3}{c}{l}&\\multicolumn{3}{c}{b}&km\\,s$^{-1}$&km\\,s$^{-1}$&&&B-M&R-S\\\\ \\\\  (1)&\\multicolumn{3}{c}{(2)}&\\multicolumn{3}{c}{(3)}&(4)&(5)&(6)&(7)&(8)&(9)\\\\ \\\\ \\hline \\\\ Abell 1060 &10$^{{\\rm h}\\hspace{-0.85ex}}$ &34$^{{\\rm m}\\hspace{-1.05ex}}$ &21\\fs 6&-27$^{\\circ\\hspace{-0.85ex}}$ &16$^{\\prime\\hspace{-0.65ex}}$ &05$^{\\prime\\prime\\hspace{-0.85ex}}$&3697 & 630 & 132 & 1 & III & C\\\\ (Hydra I)&\\multicolumn{3}{c}{269\\fdg6}&\\multicolumn{3}{c}{26\\fdg49}\\\\ \\\\ \\hline \\\\ \\end{tabular}  \\end{table*}   The observations and data reduction are described in \\S \\ref{obses}. Details of the observational technique are given in \\S \\ref{platty}, and of the process of identifying the emission-line galaxies in \\S \\ref{ident}, where a table of the detected emission-line galaxies (ELGs) is given. Measurements of radial velocities for the detected emission-line galaxies are presented in \\S \\ref{twoplate}, and those of \\mbox{H$\\alpha$+N[II]} equivalent widths and fluxes in \\S \\ref{ews}, where measured \\mbox{H$\\alpha$+N[II]} fluxes are also converted into effective star formation rates. A comparison of detected cluster emission-line galaxies in Abell 1060 with field galaxies and detected emission-line galaxies in other clusters is given in \\S \\ref{frac}. Notes on individual galaxies are given in \\S \\ref{indiv}. Finally, we present a brief discussion of our results in \\S \\ref{discus}. ",
        "conclusions": ""
    },
    "9804/astro-ph9804030_arXiv.txt": {
        "abstract": "We present the first radio observations of a sample of 13 optically and IR bright southern hemisphere classical Be stars made from the Australian Telescope Compact Array at 3.5cm and 6.3cm simultaneously. One star, $\\delta$ Cen was detected at 3.5cm and a second, $\\mu$ Cen was also thought to have been detected; further observations of this source are required to confirm this detection. No sources were detected at 6.3cm, although $\\delta$ Cen was previously detected at this wavelength by other observers at a higher flux than our detection limit. The radio observations show that the spectral energy distribution undergoes a turnover between the far IR and radio wavelengths, as was seen in previous studies.  Likewise we find no simple correlation between far IR and radio flux. Lower limits to the outer disc radius were found to be of the order of few hundred solar radii; of the order of those found previously by Taylor et al. ",
        "introduction": "Classical Be stars are defined as non supergiant B stars that have, or have had Balmer lines in emission. They are further characterised by the presence of a  continuum excess, arising from free free and bound free emission from a stellar wind. Comparison of optical and UV spectra show that two different wind regimes must coexist. Consequently  a high velocity component responsible for the high excitation lines visible in the UV, and a denser component that produces the near IR continuum excess and the optical emission lines was proposed to explain the observations. That the denser component was concentrated in the equatorial plane has long been suspected; recent  interferometric data shows that the the envelopes around Be stars are non spherical (Dougherty \\& Taylor 1992; Quirrenbach et al 1994; Stee et al 1995).  However, only a few of the brightest systems are ammeanable to an interferometric approach. Consequently, other approaches to the study of Be star circumstellar discs have been attempted. One such approach was long wavelength (mm-radio) flux measurement. When combined with near IR photometry modeling of the spectral energy distribution (SED) leads to a profile of the base density and density gradient (and hence a radial velocity law) within the circumstellar disc. However, observations of Be stars at mm and radio wavelengths showed that they were much fainter than implied by a simple extrapolation of their near IR and IRAS fluxes (e.g. Taylor et al 1990; henceforth Ta90), indicating a change in the ion density gradient within the disc. Several explanations were advanced to explain this, the most favourable being a change in disc opening angle or re-acceleration of material at large radii. Other possibilities exist, such as recombination at large radii or a truncation of the disc; see  Waters et al. (1991) and Ta90 for a thorough discussion of all the scenarios.  Given the paucity of detections in the northern hemisphere (only 6 Be stars have been detected), and the lack of multiple observations it is difficult to quantify the behaviour of the continuum excess at long wavelengths. Because of this shortfall we made observations of a sample of 13 bright southern hemisphere Be stars from the Australian Telescope Compact Array (ATCA) in 1997 April/May in an attempt to increase the sample size. ",
        "conclusions": "As a result of observations of a sample of 13 IR bright southern Be stars we have identified $\\delta$ Cen  as a radio source, which may also  be variable. A second star, $\\mu$ Cen, is  also though to be a radio emitter, although further observations are needed to confirm this. This brings the total number of Be stars detected at radio wavelengths to eight (inclusive of $\\mu$ Cen). We confirm earlier results of Ta90 that demonstrate a turnover in the spectral energy distribution between the far IR and radio wavelengths. We find no compelling evidence for a direct correlation between stellar or far IR luminosity and  radio flux. Clearly further observations of Be stars at higher sensitivities are required before such a correlation can be confirmed or rejected. Measurement of the minimum outer disc radius reveals that the circumstellar discs extend to several hundred solar radii, again reproducing the results of Ta90."
    },
    "9804/astro-ph9804340_arXiv.txt": {
        "abstract": "In the process of searching the Hubble Space Telescope archive, we have serendipitously discovered three populous Large Magellanic Cloud (LMC) clusters with ages that place them in the LMC `age gap.'. These clusters - NGC 2155, SL663, and NGC 2121 - turn out to have $[Fe/H]$$\\sim$ --1.0 and ages of $\\sim$4 Gyr. This puts them in the age gap between the intermediate-age LMC clusters, the oldest of which are $\\sim$2.5 Gyr old, and ESO121-SC03, which has an age of $\\sim$ 9 Gyr. The addition of these three clusters to the LMC age - metallicity relation has reduced the discrepancy between the age distribution of the LMC clusters and the field stars. Furthermore, it indicates that searches to find more clusters older than $\\sim$2.5 Gyr in the LMC are crucial to a better understanding of its global star formation history. ",
        "introduction": "The chemical enrichment/star formation history (SFH) is an identifying feature of every self-gravitating stellar system. One manifestation of this is the relation between age and metallicity among the star clusters in a given galaxy. Empirical information on how cluster age and metal abundance correlate provides an important clue that will eventually allow us to understand how star formation (and hence chemical enrichment) proceeds under a variety of potentially influential circumstances. It is for this reason that we strive to better define the age - metallicity relations of the cluster and field populations in galaxies. Nearby galaxies are no exception, especially the Large and Small Magellanic Clouds (L/SMC) which have provided numerous puzzles and challenges for theorists. One of the most persistent of these has been the `age gap' seen between the intermediate-age clusters ($t \\lea 2.5$ Gyr; see Sec. 4.1 for a justification of this limit) and the old clusters ($t \\gea 13$ Gyr) in the LMC (Geisler et al. \\cite{geisea1997}). The LMC age gap is also a metallicity gap in the sense that the intermediate-age clusters have $\\langle$$[Fe/H]$$\\rangle$$\\sim$$-0.5$, while the old clusters are closer to $\\langle$$[Fe/H]$$\\rangle$$\\sim$$-2$. There is only one cluster that is known to lie in the gap  - ESO121-SC03 with an age of $\\sim$9 Gyr and $[Fe/H]\\sim-1$. In contrast, recent studies based on Hubble Space Telescope (HST) data of the LMC field stars tell a signficantly different story (see Geha et al. \\cite{gehea1998} for a review). While there are variations that depend on position in the LMC (Vallenari et al. \\cite{valea1996}), the global star formation rate has been fairly constant for {\\it most} of the LMC's history. However, sometime between 2 and 4 Gyr ago, a burst of star formation occurred producing the present population of young to intermediate age field stars and clusters (Bertelli et al. \\cite{bertea1992}; Vallenari et al. \\cite{valea1996}; Gallagher et al. \\cite{gallea1996}; Holtzman et al. \\cite{holtzea1997}). Furthermore, the models produced by Geha et al. (\\cite{gehea1998}) suggest that roughly half of the LMC field stars are older than 4 Gyr. What this implies is that there should be many star clusters with ages between $\\sim$2.5 Gyr and $\\sim$13 Gyr, which were formed contemporaneously with the field stars. The obvious question of course is: where are these `age gap' clusters?  A number of investigators have undertaken surveys to find clusters in the LMC age gap (Da Costa \\cite{dac1991}; Geisler et al. \\cite{geisea1997}, and references therein). The overwhelming conclusion has been that the LMC contains only one cluster with an age between 2.5 Gyr and 13 Gyr (i.e. ESO121-SC03). There is the possibility that some clusters did exist in the age gap sometime in the past, and that these have either dissolved into the LMC field or been stripped off (Olszewski \\cite{olsz1993}). However, it is difficult to see how this process can preferentially affect only clusters formed during a given epoch, unless the mass spectrum of density fluctuations producing the gap clusters was somehow different from that which produced the other clusters in the LMC. The possible reasons for this are not obvious.   While understanding the SFH of the LMC is extremely important, this was not our original aim when we began this study. Initially, we were searching the Hubble Space Telescope (HST) archive looking for intermediate age clusters that display the red giant branch (RGB) Bump in their color-magnitude diagrams (CMD). We did find one such cluster (NGC 411; see Alves \\& Sarajedini \\cite{alsa1998}). However, more importantly for the present work, we discovered three clusters whose CMDs exhibit significant numbers of stars in the Hertzsprung gap, indicative of ages older than $\\sim$2.5 Gyr. These clusters are NGC 2155, NGC 2121, and SL663, and the next section describes the observations of these clusters and the data reduction. Section 3 presents the CMDs while Section 4 discusses the ages yielded by these CMDs. The implications of these results are detailed in Section 5. ",
        "conclusions": "The resulting relation between age and metallicity for the LMC star clusters is shown in Fig. 11. The open symbols are the ages and abundances from the work of Geisler et al. (\\cite{geisea1997}) supplemented by additional clusters from Bica et al. (\\cite{bicaea1998}) and the values for NGC 2193, Hodge 4, and ESO121-SC03 from this paper. The filled square represents the location of our three `age gap' clusters. Clearly, the clusters NGC 2155, NGC 2121, and SL 663 do indeed fall in the age gap between $\\sim$2.5 Gyr and $\\sim$13 Gyr. The reader should keep in mind, however, that if the metallicities of these clusters are higher than the values we have adopted herein, their ages will be correspondingly younger (see Sec. 4.1). As such, they may eventually be considered as belonging to the old-age tail of the IAC distribution. Future spectroscopic abundance measurements will shed more light on this.  In any event, the addition of these three clusters to the age - metallicity relation of the LMC has not eliminated the discrepancy between the cluster age distribution and that of the field stars. If there are no more clusters to be discovered in the gap, then we will require some explanation for why the clusters and the field stars exhibit such differing SFHs. However, what is {\\it more} likely is that there are as yet unstudied clusters in the LMC that will further fill in the age gap. Future ground-based and HST photometric surveys may reveal more such clusters.  For the present paper, we have utilized archival HST/WFPC2 images of LMC populous clusters to show that there are at least three clusters in the LMC age gap - NGC 2155, SL663, and NGC 2121. These clusters have $[Fe/H]$$\\sim$ --1.0 and ages of $\\sim$4 Gyr. The addition of these three clusters to the LMC age - metallicity relation is the first step in reducing the significant difference between the inferred SFHs of the LMC clusters and the field stars. This strongly indicates that searches to find more clusters older than $\\sim$2.5 Gyr in the LMC are crucial to a better understanding of its global SFH."
    },
    "9804/astro-ph9804083_arXiv.txt": {
        "abstract": "Defect models have recently been declared dead\\cite{watson97}, because they predict microwave background and matter fluctuations grossly out of line with what we see. In this talk we apply the fact that many defects are automatically destroyed at the time of radiation-matter transition, thus resurrecting the defects model. Moreover, the resurrected version predicts a cosmological constant, explains the apparent excess of hot clusters and the non-Gaussianity observed in galaxy surveys.  If this model is correct, then the MAP and PLANCK missions will not measure what people expect them to (oscillations); rather, they will measure a broad hump. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804169_arXiv.txt": {
        "abstract": "Surface brightness fluctuations (SBFs) are much brighter in the infrared than they are at optical wavelengths, making it possible to measure greater distances using IR SBFs. We report new \\Kp\\ (2.1\\micron) SBF measurements of nine galaxies in the Fornax and Eridanus clusters using a 1024$^2$-pixel HgCdTe array. We used improved analysis techniques to remove contributions to the SBFs from globular clusters and background galaxies, and we assess the relative importance of other sources of residual variance. We applied the improved methodology to our Fornax and Eridanus images and to our previously published Virgo cluster data. Apparent fluctuation magnitudes were used in conjunction with Cepheid distances to M31 and the Virgo cluster to calibrate the \\Kp\\ SBF distance scale. We find the absolute fluctuation magnitude $\\MKp\\,{=}\\,{-}5.61\\,{\\pm}\\,0.12$, with an intrinsic scatter to the calibration of 0.06 mag. No statistically significant change in \\MKp\\ is detected as a function of \\v-i. Our calibration is consistent with simple (constant age and metallicity) stellar population models. The lack of a correlation with \\v-i\\ in the context of the stellar population models implies that elliptical galaxies bluer than $\\v-i\\,{=}\\,1.2$ have SBFs dominated by younger (5--8 Gyr) populations and metallicities comparable to redder ellipticals. Significant contributions to the SBFs from anomalous populations of asymptotic giant branch stars are apparently uncommon in giant ellipticals. \\Kp\\ SBFs prove to be a reliable distance indicator as long as the residual variance from globular clusters and background galaxies is properly removed.  Also, it is important that a sufficiently high signal-to-noise ratio be achieved to allow reliable sky subtraction because residual spatial variance can bias the measurement of the SBF power spectrum. ",
        "introduction": "Because the light from a distant galaxy comes from discrete but unresolved stars, Poisson statistics lead to mottling of the galaxy's otherwise smooth surface brightness profile. Surface brightness is independent of distance, but the amplitude of the surface brightness fluctuations (SBFs) is not. As the distance ($d$) to a galaxy increases, the number of stars ($n$) in a given resolution element increases as $d^2$, but the observed flux ($f$) from each star is reduced by $d^{-2}$, making surface brightness ($nf$) independent of distance. On the other hand, the rms variation in observed flux from region to region is $n^{1/2}f$, which scales as $d^{-1}$.  The variance in surface brightness $nf^2$ normalized by the mean galaxy brightness $nf$ decreases with distance; distant galaxies appear smoother than nearby galaxies. Because the variance in surface brightness is proportional to the second moment of the stellar luminosity function, it is dominated by luminous red giant stars.  Although individual stars are not resolved, measuring SBFs probes the stellar population of the galaxy directly. With good theoretical models for stellar populations, the absolute magnitude of SBFs can be calculated, allowing a direct determination of the distance that is independent of the global dynamics or the environment of the galaxy (to the extent that the stellar populations in old stellar systems are independent of these variables). Alternatively, measuring SBFs in galaxies with known distances provides insight into stellar populations, allowing comparison with stellar evolution models and providing an empirical calibration of the SBF distance scale. Good descriptions of the theory and practice of using SBFs as a distance measurement tool and stellar population probe can be found in several papers by Tonry and collaborators (Tonry et al. 1997; Jensen, Luppino, \\& Tonry 1996, hereafter JLT; Tonry 1991; Tonry, Ajhar, \\& Luppino 1990; Tonry \\& Schneider 1988), in papers by Sodemann \\& Thomsen (1995, 1996), and in the review by Jacoby et al. (1992). The first $K$-band SBF studies are described by JLT, Luppino \\& Tonry (1993), and Pahre \\& Mould (1994).  J. Tonry and coworkers have completed an extensive survey of $I$-band SBF distances in a sample of several hundred early-type (dynamically hot elliptical and S0) galaxies which is more than 50\\% complete out to $\\sim$2800 \\kms\\ (Tonry et al. 1997). They find that the $I$-band absolute fluctuation magnitude \\MI\\ is a linear function of \\v-i\\ and has a universal zero point. Their $I$-band calibration, which is based empirically on Cepheid distances, is in good agreement with Worthey's (1993a,b; 1994) simple stellar populations models. Worthey's models predict that the effects of age and metallicity in the $I$ band are largely degenerate, so that \\MI\\ may be calibrated using a single parameter such as the \\v-i\\ color. The resulting intrinsic scatter of the $I$-band SBF distance scale is of order 0.07 mag. The purpose of the current study is to examine the behavior of the SBF calibration in the near-IR \\Kp\\ band, where Worthey's models predict a much weaker dependence on \\v-i, but potentially larger scatter as the effects of age and metallicity are no longer degenerate.  Stellar surface brightness fluctuations are very red, since they are dominated by luminous red giant stars.  The advantages of observing IR SBFs are clear:  fluctuations are ${\\sim}33$ times brighter at $K$ than at $I$, making them observable to greater distances. The SBF amplitude is inversely proportional to the seeing FWHM, and the seeing is typically much better in the IR than at optical wavelengths. Fluctuations are also red compared to the globular cluster (GC) population, so the contrast between stellar SBFs and GCs is higher at $K$. Finally, stellar population models predict that $K$-band SBF magnitudes have a much weaker dependence on \\v-i\\ than at $I$, reducing or eliminating the need to accurately measure the color of the galaxy (Worthey 1993a).  Several studies have demonstrated the feasibility of measuring IR SBFs, including Luppino \\& Tonry (1993), Pahre \\& Mould (1994), and JLT. These papers report results for a rather limited set of galaxies in the Local Group (LG) and the Virgo Cluster, but the consistency in the measured calibration of \\MK\\ is encouraging  ($\\MKp\\,{=}\\,{-}5.61\\,{\\pm}\\,0.16$ Luppino \\& Tonry, $M_{K_{sh}}\\,{=}\\,{-}5.77\\,{\\pm}\\,0.18$ Pahre \\& Mould, $\\MKp\\,{=}\\,{-}5.62\\,{\\pm}\\,0.29$ JLT). All three of these papers report absolute fluctuation magnitudes that are consistent with Worthey's (1994) predictions based on simple stellar population models. While the mean \\MKp\\ of JLT and Pahre \\& Mould (1994) are quite consistent, differences in apparent fluctuation magnitudes for individual galaxies are larger than the stated uncertainties allow. JLT discuss several possible reasons for the disagreement in fluctuation magnitudes and showed that residual variance resulting from spatial variations in the detector sensitivity and dark current can contribute significantly to the power spectrum in low signal-to-noise ratio ($S/N$) observations. A residual pattern of only 0.1\\% of the sky level can significantly change the fluctuation magnitude measured. To use IR SBFs as a reliable distance indicator, it is critical that the observations be of sufficiently high $S/N$ ratio to avoid biases from residual variances. JLT also demonstrated that properly subtracting the sky and carefully sampling the point spread function (PSF) are crucial to accurately measuring the fluctuation amplitude. An uncertainty in the background subtraction of 2\\% of the sky level can dominate the uncertainty in the fluctuation magnitude.  Our Virgo sample showed a large dispersion (0.29 mag) due to the depth of the cluster and observational errors.  Pahre \\& Mould (1994) also addressed the dispersion in fluctuation magnitude in a similarly-sized sample of Virgo ellipticals. Clearly a larger sample of galaxies is needed to calibrate the IR SBF distance scale.  To better understand the effects stellar populations have on the fluctuation amplitude, we observed five galaxies in the Fornax cluster.  Fornax is ideal for several reasons: first, it is much more centrally concentrated than the Virgo cluster. The reduced dispersion in distances allows us to better quantify and understand the uncertainty in our SBF measurements. Second, Fornax contains a large number of giant elliptical galaxies with a wide range of \\v-i\\ colors, metallicities, and globular cluster populations.  The Virgo galaxies observed by JLT and Pahre \\& Mould (1994) spanned a very limited range in \\Mg2\\ index to avoid stellar population variations which may affect the calibration of \\MK. We now extend the sample to calibrate \\Kp\\ fluctuation magnitudes across a wider range of metallicities and \\v-i\\ colors to compare with theoretical stellar population models. Finally, Cepheid distances to several Virgo cluster galaxies and to the Fornax spiral NGC~1365 have been measured using the Hubble Space Telescope (HST). We can empirically anchor our \\Kp-band SBF calibration both to the Cepheid and the $I$-band SBF distance scales.  This paper reports results from our observations of four Eridanus cluster elliptical galaxies in addition to the Fornax galaxies. Eridanus is not as compact a cluster as Fornax, but we can still compare our IR SBF results with the $I$-band results from Tonry (1997). Finally, we return to our SBF measurements for the seven Virgo ellipticals discussed by JLT and apply the improved techniques we describe in this paper.  We present the updated results for these galaxies in Section~\\ref{virgorevisited}, along with new observations of NGC~4365. ",
        "conclusions": "  $K$-band SBFs can be measured reliably and used to determine distances to early-type galaxies.  Because SBFs are relatively bright in the IR and the seeing is typically very good, integration times can be quite modest.  However, observations must be sufficiently deep to adequately sample the GC luminosity function and remove sources of residual variance.  Low-$S/N$ ratio measurements are unreliable. Optical images can be used to identify and remove GCs and background galaxies that are not detected in the IR, improving the IR SBF measurement. We empirically calibrated the \\Kp\\ SBF distance scale using Cepheid distances to M31 and the Virgo cluster. The absolute \\Kp\\ fluctuation magnitude is $\\MKp\\,{=}\\,{-}5.61\\,{\\pm}\\,0.06$ (statistical error) with a total uncertainty of 0.12 mag. No significant change in \\MKp\\ with \\v-i\\ was observed over the range in color spanned by the galaxies in this sample. Accurate color measurements are not required to measure the \\Kp\\ SBF distance to a galaxy.  \\Kp\\ SBF magnitudes are consistent with predictions from simple stellar population models. The lack of a correlation in \\MKp\\ with \\v-i\\ is best explained by a spread in ages among the galaxies observed.  The redder ellipticals are consistent with 12 to 17 Gyr stellar population models, while the bluer galaxies in our sample must have younger 5 to 8 Gyr populations. Metallicities appear to vary less than ages, with the typical galaxy having [Fe/H]\\,=\\,${-}0.25$. The stellar population models show that the age-metallicity degeneracy is broken with $K$-band SBFs, allowing one to distinguish between old, metal-poor and young, metal-rich populations. Examining the radial variation in IR SBFs will help distinguish between different galaxy formation scenarios. Our observations do not agree with the relationship between \\MKp\\ and the \\Mg2\\ index predicted by the stellar population models."
    },
    "9804/astro-ph9804263_arXiv.txt": {
        "abstract": "The high redshift radiogalaxy 1243+036 ($z=3.6$) presents an asymmetric \\lya\\ profile of FWHM 1550 \\kms\\ as measured by van~Ojik \\etal\\ We propose that the blue asymmetry in the \\lya\\ profile is not due to narrow absorption dips but consists of narrow emission peaks. We interpret the blueshifted peaks near $-1130, -850$ and $-550 \\kms$ (relative to the peak of full profile) as being the result of Fermi acceleration of \\lya\\ produced by jet-induced star formation in the wake of a 300~\\kms\\ shock. This shock would be caused by the deflection of the radio-jet at the observed position of the radio bend which also coincides spatially with the excess \\lya\\ emission reported by van~Ojik \\etal\\ ",
        "introduction": "\\label{intro}  Lyman $\\alpha$ is the strongest line observed in very high redshift radio galaxies (HZRG). Although the brightness of \\lya\\ peaks near or at the nuclear position, most of the emission is spatially resolved with the fainter emission extending up to radii 40--130\\,kpc. The \\lya\\ profile in HZRG is characterized by a FWHM in the range 700--1600\\kms\\ (van~Ojik 1995 and references therein).  The intermediate resolution study of \\lya\\ profiles carried out by van~Ojik \\etal\\ (1997) has revealed the presence of troughs which are well explained by \\hi\\ gas absorption present in the environment of the parent radio galaxy. In their sample, however, the HZRG 1243+036 ($z=3.6$) may call for a different interpretation, namely that the broad profile presents true narrow emission features in between the narrow `dips' which van~Ojik \\etal\\ (1996: vO96) have interpreted as absorption features. The repeated scattering of the resonance \\lya\\ line across a shock discontinuity was shown by Neufeld \\& McKee (1988: NM88) to result in a systematic blueshift of \\lya.  In this paper, we develop in more detail the Fermi acceleration model and propose that the narrow features observed by vO96 correspond to a small number of across-shock scatterings. ",
        "conclusions": ""
    },
    "9804/astro-ph9804325_arXiv.txt": {
        "abstract": "Recent discoveries by the {\\it Rossi X-Ray Timing Explorer} indicate that most of the rapidly accreting ($\\dot M \\gtrsim 10^{-11} M_\\odot \\ {\\rm yr}^{-1}$) weakly magnetic ($B\\ll 10^{11} \\ {\\rm G}$) neutron stars in the Galaxy are rotating at spin frequencies $\\nu_s \\gtrsim 250 \\ {\\rm Hz}$. Remarkably, they all rotate in a narrow range of frequencies (no more than a factor of two, with many within 20\\% of 300 Hz). I suggest that these stars rotate fast enough so that, on average, the angular momentum added by accretion is lost to gravitational radiation. The strong $\\nu_s$ dependence of the angular momentum loss rate from gravitational radiation then provides a natural reason for similar spin frequencies. Provided that the interior temperature has a large scale asymmetry misaligned from the spin axis, then the temperature sensitive electron captures in the deep crust can provide the quadrupole needed ($\\sim 10^{-7} M R^2$) to reach this limiting situation at $\\nu_s\\approx 300$ Hz. This quadrupole is only present during accretion and makes it difficult to form radio pulsars with $\\nu_s>(600-800) \\ {\\rm Hz}$ by accreting at $\\dot M \\gtrsim 10^{-10} M_\\odot \\ {\\rm yr^{-1}}$.  The gravity wave strength is $h_c\\sim (0.5-1) \\times 10^{-26}$ from many of these neutron stars and $>2\\times 10^{-26}$ for Sco X-1. Prior knowledge of the position, spin frequency and orbital periods will allow for deep searches for these periodic signals with gravitational wave interferometers (LIGO, VIRGO and the ``dual-recycled'' GEO 600 detector) and experimenters need to take such sources into account. Sco X-1 will most likely be detected first. ",
        "introduction": "The launch of the {\\it Rossi X-Ray Timing Explorer} (RXTE) has allowed for the discovery of fast quasi-periodic variability from many rapidly accreting ($ \\dot M\\gtrsim 10^{-11} M_\\odot \\ {\\rm yr^{-1}}$) neutron stars. These observations strongly suggest that these neutron stars (NSs) are rapidly rotating, as predicted by those scenarios connecting the millisecond radio pulsars to this accreting population (see Bhattacharya 1995 for an overview). Strohmayer et al. (1996) were the first to detect nearly coherent $\\nu_B=363$ Hz oscillations during type I X-ray bursts from the low accretion rate ($\\dot M< 10^{-9} M_\\odot \\ {\\rm yr^{-1}}$) NS 4U~1728-34. Pulsations were detected in six of the eight bursts analyzed at that time. In addition, two high frequency quasi-periodic oscillations (QPOs) were seen in the persistent emission. These changed with accretion rate, but maintained a fixed difference frequency of $\\nu_d\\approx 363 $ Hz, identical to the period seen during the bursts.  The detection of two drifting QPO's (in the persistent emission) separated by a fixed frequency identical to that seen in the bursts naturally leads to beat frequency models (Strohmayer et al. 1996; Miller, Lamb, \\& Psaltis 1998). The difference frequency is presumed to be the NS spin frequency, $\\nu_s$, whereas the upper frequency has different origins in different models (see van der Klis 1998 for a summary). In addition, the temporal behavior of the periodic oscillations both during the rise of the bursts (Strohmayer, Zhang, \\& Swank 1997b) and in the cooling tails (Strohmayer et al. 1997a) are most easily explained in terms of rotation.  There are six NSs with measured periodicities during Type I X-ray bursts (see Table 1). Both the difference frequencies ($\\nu_d$) and the burst frequencies ($\\nu_B$) are in a narrow range, from 260 to 589 Hz. For two objects (KS 1731-260 and 4U 1636-53) the difference frequencies are one-half the burst values. Which value is $\\nu_s$ is not resolved.  There are also many NSs that accrete at higher rates and are not regular Type I X-ray bursters. Many of these objects, notably the ``Z'' sources, also show drifting QPO's at fixed separation, again with a similarly narrow frequency range (roughly 250-350 Hz). Beat-Frequency like models are also applied to these observations so as to infer $\\nu_s$. The applicability of such a model is less clear when the difference frequency is not constant (Sco X-1, van der Klis et al. 1997; 4U 1608-52, Mendez et al. 1998).  If accreting matter always arrives with the specific angular momentum of a particle orbiting at the NS radius ($R=10 R_6 {\\rm km}$), then it only takes $\\sim 10^7 \\ {\\rm yrs}$ of accretion at $\\dot M\\approx 10^{-9} M_\\odot \\ {\\rm yr^{-1}}$ for a $M=1.4M_{1.4} M_\\odot$ NS to reach $\\nu_s=50$ Hz from an initially low frequency. It is thus remarkable that these NSs are all rotating at nearly the same rate. White and Zhang (1997) argued that this similarity arises because these NSs are magnetic and have reached an equilibrium where the magnetospheric radius equals the co-rotation radius. This requires an intrinsic relation between their magnetic dipoles, $\\mu_b$, and $\\dot M$ so that they all reach the same rotational equilibrium (most likely $\\mu_b\\propto \\dot M^{1/2}$) and a way of hiding the persistent pulse typically seen from a magnetic accretor.  My alternative explanation for these spin similarities is that gravitational wave (GW) emission has started to play an important dnrole. If the NS has a misaligned quadrupole moment, $Q$, then the strong spin frequency dependence of GW emission defines a critical frequency beyond which accretion can no longer spin-up the star.  Such a NS will radiate energy via GW's at the rate $\\dot E=32 GQ^2\\omega^6/5 c^5$, where $\\omega=2\\pi \\nu_s$, and lose angular momentum at the rate $N_{gw}=\\dot E/\\omega$. Balancing this spin-down torque with the characteristic spin-up torque from time-averaged accretion, $N_a\\approx \\dot M(GMR)^{1/2}$, gives the $Q$ needed so as to make the critical frequency 300 Hz, \\begin{equation}\\label{eq:qneed} Q\\approx 4.5 \\times 10^{37} \\ {\\rm g \\ cm^2}\\left(\\dot M\\over 10^{-9} \\ {\\rm M_\\odot \\ yr^{-1}}\\right)^{1/2}\\left(300 \\ {\\rm Hz}\\over \\nu_s\\right)^{5/2}, \\end{equation} or $<10^{-7}$ of the NS moment of inertia, $I\\approx 10^{45} \\ {\\rm g \\ cm^2}$. The similarities in $\\nu_s$ may then arise because of the weak dependencies of the critical frequency on $Q$ and $\\dot M$.  What is the source of the misaligned quadrupole?  Wagoner (1984) argued that accreting NSs would get hung-up at spin frequencies where the Chandrasekhar-Friedman-Schutz (CFS) instability sets in. However, Lindblom (1995) and Lindblom \\& Mendell (1995) have shown that the star needs to be very near the breakup frequency ($\\nu_s \\gtrsim \\ {\\rm kHz}$) for such an instability to occur, even for the core temperatures $T_c=(1-3)\\times 10^8 \\ {\\rm K}$ of rapidly accreting NSs (Ayasli \\& Joss 1978; Brown \\& Bildsten 1998). The spin frequencies for these NSs are too slow for such an instability.  I present in \\S 2 a new source for lateral density variations in an accreting NS; electron captures (hereafter EC) in the crust. The constant compression of the crust forces nuclei to undergo EC when the electron Fermi energy, $E_F$, is high enough to make a transition. However, the crust is hot enough in a rapidly accreting NSs to make the EC rates temperature sensitive. Hotter regions then capture at lower pressures, so that the density jump associated with the EC is at a higher altitude in the hotter parts of the crust. Moderate lateral temperature variations then lead to density variations large enough to generate the required $Q$. This outcome is independent of the particular source of the temperature variations. One possible cause for a $T$ asymmetry relative to the spin axis is a weak magnetic field. I conclude in \\S 3 by finding the GW signal strength and estimating detection. ",
        "conclusions": ""
    },
    "9804/astro-ph9804113_arXiv.txt": {
        "abstract": "\\vspace{-3ex} We present a numerical simulation of the dynamical collapse of a nonrotating, magnetic molecular cloud core and follow the core's evolution through the formation of a central point mass and its subsequent growth to a $1~\\msol$ protostar. The epoch of point-mass formation  (PMF) is investigated by a self-consistent extension of previously presented models of core formation and contraction in axisymmetric, self-gravitating, isothermal, magnetically supported interstellar molecular clouds. Prior to PMF, the core is dynamically contracting and is not well approximated by a quasistatic equilibrium model. Ambipolar diffusion, which plays a key role in the early evolution of the core, is unimportant during the dynamical pre-PMF collapse phase. However, the appearance of a central mass, through its effect on the gravitational field in the inner core regions, leads to a ``revitalization'' of ambipolar diffusion in the weakly ionized gas surrounding the central protostar. This process is so efficient that it leads to a decoupling of the field from the matter and results in an outward-propagating hydromagnetic C-type shock. The existence of an ambipolar diffusion-mediated shock of this type was predicted by Li \\& McKee (1996), and we find that the basic shock structure given by their analytic model is well reproduced by our more accurate numerical results. Our calculation also demonstrates that ambipolar diffusion, rather than Ohmic diffusivity operating in the innermost core region, is the main field decoupling mechanism responsible for driving the shock after PMF.  The passage of the shock leads to a substantial redistribution, by ambipolar diffusion but possibly also by magnetic interchange, of the mass contained within the magnetic flux tubes in the inner core. In particular, ambipolar diffusion reduces the flux initially threading a collapsing $\\sim 1~\\msol$ core by a factor $\\simgt 10^3$ by the time this mass accumulates within the inner radius ($\\simeq 7.3~{\\rm AU}$) of our computational grid. This reduction, which occurs primarily during the post-PMF phase of the collapse, represents a significant step towards the resolution of the protostellar magnetic flux problem.  Our calculations indicate that a $1~\\msol$ protostar forms in $\\sim 1.5 \\times 10^5~\\rm{yr}$ for typical cloud parameters. The mass accretion rate is time dependent, in part because of the C-shock that decelerates the infalling matter as it propagates outward: the accretion rate rises to $\\simeq 9.4~\\msol~\\rm{Myr}^{-1}$ early on and decreases to $\\simeq 5.6~\\msol~{\\rm{Myr}}^{-1}$ by the time a solar-mass protostar is formed. The infalling gas disk surrounding the protostar has a mass $\\sim 10^{-2}~\\msol$ at radii $r \\simgt 500~\\rm{AU}$. A distinguishing prediction of our model is that the rapid ambipolar diffusion after the formation of a protostar should give rise to large ($\\simgt 1~{\\rm{km}}~{\\rm s}^{-1}$), and potentially measurable, ion--neutral drift speeds on scales $r \\simlt 200~\\rm{AU}$.  The main features of our simulation, including the C-shock formation after PMF, are captured by a similarity solution that incorporates the effects of ambipolar diffusion (Contopoulos, Ciolek, \\& K\\\"{o}nigl 1997). ",
        "introduction": "It is generally accepted that most of the star-formation activity in our galaxy takes place through the gravitational collapse of molecular cloud cores (e.g., Mouschovias 1987; Shu, Adams, \\& Lizano 1987). It is, furthermore, believed that interstellar magnetic fields play a central role in this process in that their stresses are the dominant agent that acts against gravity to prevent, or delay, cloud collapse (e.g., Mouschovias 1978; McKee et al. 1993). This is embodied in the concept of a critical mass $M_{\\rm crit}$, which in general takes account of both the magnetic and the thermal pressure contributions to the support of the cloud, but which, in the case that magnetic stresses dominate, reduces to $M_{\\rm crit} \\approx 0.13 \\phiB/G^{1/2}$, where $G$ is the gravitational constant and $\\phiB$ is the magnetic flux that threads the cloud. Clouds whose mass $M$ exceeds $M_{\\rm crit}$ are ``supercritical'': they collapse on the free-fall timescale. In contrast, ``subcritical'' clouds ($M < M_{\\rm crit}$) can avoid collapse on the much longer ambipolar diffusion timescale. In the latter case, the neutrals gradually contract by diffusing inward through the ions and field, leaving behind a magnetically supported envelope and eventually forming a supercritical core that undergoes dynamical collapse (e.g., Mouschovias 1996).  Because of the complexity of the problem --- it involves solving the full dynamical equations of a magnetized, multicomponent (neutrals, ions, electrons, as well as charged and neutral grains) fluid that evolves over many decades in size and density in a nonspherically symmetric manner (because of the presence of ordered magnetic fields and likely also rotation) --- much of the progress in this area has been accomplished through the use of numerical simulations. One of the main efforts to simulate core formation and contraction due to ambipolar diffusion in magnetically supported molecular clouds has been carried out by Mouschovias and coworkers (e.g., Fiedler \\& Mouschovias 1992, 1993; Ciolek \\& Mouschovias 1993, 1994, 1995, hereafter CM93, CM94, CM95, respectively; Basu \\& Mouschovias 1994, 1995a, b). These studies followed the evolution of the core over six decades in density up to central densities $\\sim 3 \\times 10^9~\\cc$, where the assumption of isothermality starts to break down because of radiative trapping (e.g., Gaustad 1963; Hayashi 1966). This assumption had been adopted in the interest of simplicity: sophisticated and computationally intensive numerical techniques are generally needed to calculate the thermal structure of the gas during the opaque phase of protostellar evolution (e.g., Larson 1969, 1972; Winkler \\& Newman 1980; Stahler, Shu, \\& Taam 1981; Boss 1984; Myhill \\& Kaula 1992; Myhill \\& Boss 1993). As a result of this restriction, the aforementioned calculations did not follow the collapse of the core to the time where a point mass -- a protostar -- is formed at the center, although they did obtain valuable information on the conditions leading to this critical event. In particular, by the time these simulations were terminated, the inner region of the core was collapsing dynamically and was characterized by neutral infall speeds $\\sim C$ (the isothermal speed of sound) and inward accelerations $\\simgt 0.3 |\\gr|$ [where $\\gr(r)$ is the local gravitational acceleration]. Furthermore, the thermal pressure, while remaining relatively unimportant in the envelope, came to exceed the magnetic pressure near the center. Basu (1997) derived a time-dependent, semianalytic solution that extended these ambipolar diffusion models up to the instant of point-mass formation (hereafter referred to as PMF \\footnote{Creation of a central point mass was commonly referred to in previous papers as ``core formation.'' In this paper we use the term ``point-mass formation'' so as not to confuse this process with the formation of an extended, magnetically supercritical, molecular cloud core.}). He found that ambipolar diffusion continues to gradually erode the retarding magnetic forces in the inner core, making the collapse increasingly more dynamical (and the thermal-to-magnetic pressure ratio in the inner core progressively larger) as PMF is approached.  The diminution of magnetic forces in the innermost regions of a collapsing core just prior to PMF suggests that one could gain some insight into the protostar formation process from previous studies of PMF in {\\it nonmagnetic}, spherically symmetric, isothermal clouds.  Analytic similarity solutions have uncovered two limiting behaviors: Penston (1969) and Larson (1969) found a solution in which, just before PMF, the infall speed approaches $\\sim 3.3~C$ at all radii while the density scales with radius as $r^{-2}$, resulting in a spatially uniform mass inflow rate $\\sim 29~C^3/G$ (where $G$ is the gravitational constant). Hunter (1977) extended this solution past PMF and showed that, immediately after the central mass is formed, the accretion rate onto the protostar increases to $\\sim 47~C^3/G$. In the other limit, Shu (1977) obtained a solution that is static prior to PMF (with the density distribution of a singular isothermal sphere, which also scales as $r^{-2}$) and that takes on an expansion-wave character (with a constant mass accretion rate $\\sim 0.98~C^3/G$ onto the protostar) following PMF. Hunter (1977) and Whitworth \\& Summers (1985) demonstrated that there are, in fact, infinitely many similarity solutions that span the range between the Larson-Penston and Shu results, with the nature of any given solution being determined by the initial configuration of the cloud and the conditions at its boundary. Numerical simulations carried out by Hunter (1977) and by Foster \\& Chevalier (1993) confirmed the dependence on the initial and boundary conditions. In particular, it was found that the behavior of the central regions of clouds that are initially marginally stable to collapse approximates that of the Larson-Penston solution at the PMF epoch, although it was determined that the mass accretion rate onto the protostar declines at later times. It was, however, also found that the post-PMF evolution of clouds that initially have more extended envelopes approximates that of the Shu solution at late times. Since the initial and boundary conditions of real clouds are expected to depend on the detailed configuration and evolution of the embedded magnetic field, it is clear that one needs to incorporate magnetic field effects into the collapse calculations to adequately model the formation of protostars.  There have been several recent attempts to calculate PMF following the collapse of magnetic interstellar clouds. Although they have all contributed to our understanding of the processes involved, their results were hampered by the adopted  assumptions or approximations. For example, Tomisaka (1996) modeled clouds that had equal thermal and magnetic energy densities, so that they were not primarily supported by magnetic fields. This means that his model clouds were magnetically supercritical. This assumption is at variance with H {\\sc I} and OH Zeeman measurements of magnetic field strengths in molecular clouds, which are consistent with models of magnetically subcritical clouds (Crutcher et al. 1993, 1994, 1996). Li \\& Shu (1997) modeled PMF in self-similar, magnetic cores. They assumed that cores immediately before PMF can be represented by hydrostatic configurations of singular isothermal disks and that the magnetic flux is frozen into the neutrals. These assumptions are inconsistent with the above-cited results of numerical simulations and semianalytic solutions of the collapse of magnetically supported molecular clouds that undergo ambipolar diffusion (as well as with the simulations of thermally supported spherical clouds), which have found that the inner core regions collapse dynamically as PMF is approached (see also \\S 3.2). As we show below, ambipolar diffusion, which plays a key role in bringing about the dynamical collapse, is generally important also {\\it after} PMF. Safier, McKee, \\& Stahler (1997) did include ambipolar diffusion in the modeling of magnetic cloud collapse and post-PMF evolution. However, their model is strictly spherical, and they neglected the effect of thermal pressure gradients in comparison with magnetic stresses. Furthermore, their formulation is quasistatic and does not incorporate the magnetic induction equation for the time evolution of the magnetic field. Li (1998) extended the Safier et al. model by including thermal-pressure and time-dependent terms and by adding the induction equation. This enabled him to follow the time evolution of his model cores (for $r > 150~\\rm{AU}$) even during the dynamical phases of the collapse. However, by retaining the spherical-symmetry assumption of Safier et al., his calculations were also unable to yield the geometry of the magnetic field lines. \\footnote{The same is true for the spherically symmetric self-similar model of a collapsing magnetic cloud devised by Chiueh \\& Chou (1994), which, however, does not include ambipolar diffusion. It should be noted that all models that assume spherical symmetry do not satisfy the solenoidal condition $\\nabla \\cdot \\Bvec =0$ on the magnetic field.}  The importance of ambipolar diffusion after PMF can be inferred by comparing the ambipolar diffusion timescale $\\tad = r/\\vd$ (where $\\vd$ is the ion--neutral drift speed) and the gravitational contraction ($\\simeq$ free-fall) timescale $\\tgr = (r/|\\gr|)^{1/2}$ before and after PMF. In axisymmetric geometry, $\\tad/\\tgr \\simeq (\\tgr/\\tni) \\, {\\mu_ B}^2$ in the inner flux tubes of a supercritical core (e.g., Mouschovias 1991), where ${\\mu_B}(r)$ is the total mass-to-flux ratio at radius $r$ (in units of the critical value for gravitational collapse) and \\begin{equation} \\label{tnieq} \\tni = 1.4 \\left[1 + 0.067 \\frac{\\left(\\mHII /2~\\rm{a.m.u.}\\right)}{\\left(\\mi/30~\\rm{a.m.u.}\\right)}\\right] \\frac{1}{\\nni \\sigw} \\end{equation} is the neutral--ion collision time. In equation (\\ref{tnieq}), $\\nni$ is the ion density and $\\sigw$ is the average collisional rate between ions of mass $\\mi$ and neutral $\\HII$ molecules of mass $\\mHII$ ($\\simeq 1.7 \\times 10^{-9}~{\\rm{cm^3}}~{\\rm s}^{-1}$ for collisions between neutrals and $\\Mgp$ or $\\HCOp$ ions; McDaniel \\& Mason 1973); the factor 1.4 accounts for a 20\\% helium abundance by number. CM94 and CM95 found that, to first order for the late pre-PMF evolution of cores in disk-like clouds, the magnetic field $B \\approx 3\\, (r_0/r)~\\rm{mG}$ (where $r_0 = 40~\\rm{AU}$), $\\nni \\approx 0.1~\\cc$ (valid for neutral densities $\\nn \\simgt 10^7~\\cc$; see Figs. $2b$ and $4b$ in CM94), the vertical column density $\\sign \\approx 5\\, (r_0/r)~{\\rm g}~\\rm{cm}^{-2}$, and the total mass $M(r) \\approx 6 \\times 10^{-3} (r/r_0) \\msol$. For a disk-like cloud, the critical mass-to-flux ratio $(M/\\phiB)_{\\rm{d,crit}}= (4 \\pi^2 G)^{-1/2}$, where $G$ is the gravitational constant (Nakano \\& Nakamura 1978). Therefore, one finds $\\mu_B \\approx 2.7$, $|\\gr| \\approx G M(r)/r^2 \\approx 2 \\times 10^{-6} (r_0/r) ~{\\rm{cm}}~\\rm{s}^{-2}$, and $\\tni \\approx 7 \\times 10^9~\\rm{s}$, which yields $\\tgr \\approx 2 \\times 10^{10} (r/r_0) ~\\rm{s}$ and $\\tad/\\tgr \\approx 20\\, (r/r_0)$. It follows that $\\tad/\\tgr \\gg 1$ for $r \\gg 2~\\rm{AU}$. Hence, as has already been demonstrated in the past, ambipolar diffusion is ineffective as a dynamically collapsing core approaches PMF, and for $r \\gg 2~\\rm{AU}$ the magnetic flux can be considered frozen into the neutrals. (For reasons discussed in \\S 2.2, we do not consider $r \\simlt 5~\\rm{AU}$.) Turning now to the post-PMF epoch, when the central point mass comes to dominate the gravitational field in the innermost flux tubes, we use the relation $\\tad/\\tgr \\approx \\left(1- |\\an|/|\\gr|\\right)^{-1} (\\tgr/\\tni)$, where $\\an$ is the inward acceleration of the neutrals and $\\gr$ is the total gravitational acceleration (see eqs. [\\ref{rforceeq}] and [\\ref{tadeq}] in \\S 3.3). In this case, normalizing $r$ as before, $|\\gr| \\approx 4 \\times 10^{-4} (\\mcent/\\msol) (r_0/r)^2 {\\rm{cm}}~\\rm{s}^{-2}$. Substituting again $\\tni \\approx 7 \\times 10^9~{\\rm s}$ (corresponding to $\\nni \\approx 0.1~\\cc$) , the above ratio becomes $\\tad/\\tgr \\approx 0.2\\, (r/r_0)^{3/2}(\\msol/\\mcent)^{1/2}\\left(1-|\\an|/|\\gr|\\right)^{-1}$. For $|\\an|/|\\gr|$ in the range 0.2 -- 0.9, which corresponds to the period after PMF when the collapse becomes progressively more dynamical, one infers $\\tad/\\tgr \\approx (0.2 - 2) (r/r_0)^{3/2}(\\msol/\\mcent)^{1/2}$. Hence $\\tad \\simlt \\tgr$ for $\\mcent \\simgt 0.1 \\msol$ and $r \\simlt 50~\\rm{AU}$. [A similar result is obtained if one continues to use the pre-PMF relations and simply substitutes $\\mcent$ for $M(r)$.] This estimate indicates that {\\em decoupling of the gas and magnetic field by ambipolar diffusion should occur in the inner core regions after PMF}. The physical reason for this is that the strength of the gravitational field in the weakly ionized gas near the origin is greatly enhanced by the appearance and growth of a central point mass, causing the neutrals to fall in more rapidly while the plasma and magnetic field are left behind. (The same basic reason --- the appearance of a progressively growing free-fall zone around the origin following PMF --- is also the cause of the increase in the mass inflow rate into the center at that epoch first discovered in the above-referenced nonmagnetic collapse calculations.) The foregoing conclusion is verified by a detailed calculation in \\S~3.3, where we show that ambipolar diffusion after PMF is, in fact, so efficient that it effectively decouples the neutrals and magnetic field in the innermost core region, with dramatic consequences for the subsequent dynamical evolution of the core. An alternative, yet equivalent, analysis of the effectiveness of ambipolar diffusion following PMF, based on the scaling of magnetic forces (particularly the magnetic tension force) after PMF, is given in Appendix C.  In this paper we present a detailed numerical simulation of point-mass formation in a nonrotating, magnetic, dynamically collapsing protostellar core, properly accounting for the effect of ambipolar diffusion. Unlike earlier studies, we use an initial state that is consistent with the realistic models of core formation and collapse, as presented earlier by Mouschovias and coworkers. As we noted above, the simulations carried out by that group were terminated at densities where the isothermality assumption started to become invalid because of radiative trapping. Although a proper treatment of radiative trapping is indispensable for a complete treatment of the star-formation process, one can adopt a simpler approach that circumvents this difficulty by removing the region of radiative trapping from the active computation mesh. This is justified by the fact that the region of radiative trapping (typically a few AU) is several orders of magnitude smaller than the characteristic core size ($r_{\\rm{core}} \\approx 0.1~\\rm{pc}$). This region can therefore be considered to be effectively point-like, and one can proceed to calculate the formation of a central point mass within a gravitationally collapsing core and its effect on the subsequent evolution of the core by retaining the isothermality assumption. In adopting this approach, we note that the assumption of isothermality was also employed in previous studies of protostar formation in nonmagnetic cloud cores as well as in the more recent attempts to model PMF in magnetic clouds.  The plan of the paper is as follows. In \\S~2 we review the main characteristics of the pre-PMF core-evolution calculations of CM93, CM94, and CM95 and outline the model modifications that we have implemented to extend the simulations beyond PMF. In \\S~3 we present the results of our calculations. We consider a typical model and follow the evolution of all physical quantities of interest through the PMF epoch. We also describe the formation (due to ambipolar diffusion in the innermost core after PMF) of a C-type shock and consider its propagation through the collapsing core. The appearance of such a shock as a result of field--matter decoupling was first pointed out by Li \\& McKee (1996), who proposed that the relevant field decoupling mechanism was Ohmic dissipation in the innermost regions of the core. Our study, however, reveals that ambipolar diffusion occurring outside the region of Ohmic dissipation is the main cause of magnetic flux decoupling after PMF. In \\S~4 we discuss the structure of the ambipolar diffusion-mediated shock and show that our numerical calculations qualitatively reproduce the predictions of the simplified analytic model constructed by Li \\& McKee (1996). We present a quantitative comparison with their results and also address the issue of interchange instability in the post-shock region, first raised in their work, in light of our detailed computations. In that section we also discuss the observational implications of our simulation and briefly comment on the magnetic flux problem during star formation. Our results are summarized in \\S~5. ",
        "conclusions": "\\subsection{Observational Comparisons and Predictions} We may compare our result for the protostellar accretion rate during the PMF epoch with observations of star-forming cores. As shown in \\S~3.3, the accretion rate rises rapidly early on to $\\mcentdot \\simeq 9.4~\\msol~\\rm{Myr}^{-1}$ for $\\delt \\simlt 10^3~\\rm{yr}$ (see Fig. $2b$) and stays at this value up to the formation of the hydromagnetic shock. For $\\delt \\simgt 4 \\times 10^3~\\rm{yr}$ the shock is able to decelerate the infalling matter, and the accretion rate decreases to $\\mcentdot \\simeq 5.6~\\msol~\\rm{Myr}^{-1}$ by $\\delt \\simeq 1.5 \\times 10^5~\\rm{yr}$ (the time when $\\mcent = 1~\\msol$). Therefore $\\mcentdot$ decreases with increasing central mass (see Fig. $2c$). This is consistent with estimates of ages and accretion rates ($\\propto t_{\\rm{age}}^{-1}$) for young stellar objects, as deduced from evolutionary diagrams inferred from observations of Class 0 and Class I objects (e.g., Saraceno et al. 1996). In particular, the lifetimes of Class 0 objects were estimated in this way to be an order of magnitude shorter than those of Class I objects, providing evidence for a decrease in the protostellar accretion rate as an object evolves from a Class 0 source to a Class I source (e.g., Andr\\'{e} 1995; Ward-Thompson 1996). Another argument for a time-dependent accretion rate was given by Bontemps et al. (1996), who analyzed the observed CO momentum flux of several young stellar objects and found a noticeable decline in the CO flux with decreasing circumstellar envelope mass. They suggested that this is indicative of a decrease in the protostellar accretion rate (which they assumed to be proportional to the mass outflow rate) as an object evolves from Class 0 to Class I.  It is also of interest to note the ${}^{13}{\\rm{CO}}(J=1-0)$ observations of infalling disks for the protostellar candidates HL Tauri (Hayashi, Ohashi, \\& Miyama 1993) and L1551-IRS5 (Ohashi et al. 1996). HL Tauri has a mass $\\sim 0.6~\\msol$ and a surrounding disk with radius $\\sim 1400~\\rm{AU}$ and mass $\\sim 0.03~\\msol$. From the observed kinematics Hayashi et al. derive an accretion rate $\\sim 9~\\msol~\\rm{Myr}^{-1}$ at $r \\sim 700~\\rm{AU}$. The embedded protostar L1551-IRS5 has a mass $\\sim 0.5~\\msol$ and is surrounded by a disk with radius $\\sim 700~\\rm{AU}$ and mass in the range $3.9 \\times 10^{-2} - 8.1 \\times 10^{-2}~\\msol$. Ohashi et al. deduce an accretion rate in the range $13 - 26~\\msol~{\\rm{Myr}}^{-1}$ at $r \\sim 600~\\rm{AU}$. These values are comparable to our model results for $\\delt_5 \\simlt \\delt \\simlt \\delt_6$ (corresponding to $0.2~\\msol \\simlt \\mcent \\simlt 1~\\msol$; see Fig. $6a$). During this period, $ 6~ \\msol~{\\rm{Myr}}^{-1} \\simlt \\mdot \\simlt 9~ \\msol~{\\rm{Myr}}^{-1}$ for $r \\simgt 500~\\rm{AU}$ (see Fig. $6g$). [Note, however, that the temperatures of HL Tauri and L1551-IRS5 are in the range $15 - 50~{\\rm K}$, which is greater than our assumed value of 10 K and should lead to higher values of $\\mdot$ and $\\mcent(t)$; e.g., Shu et al. 1987. For a discussion of how quantities scale with temperature in our models, see Basu \\& Mouschovias 1994.] In Figure 7 we show the mass ($M - \\mcent$) of the gas surrounding the central point mass in our typical model as a function of $r/\\rzero$ for the same seven times $\\delt_j$ as in Figure 6. (Taken together, Figs. $6a$ and 7 may be taken to represent the evolution of a protostar from a Class 0 to a Class I object.) For times $\\simgt \\delt_5$, the surrounding disk mass spans the range $0.01 - 0.1~\\msol$ for $r \\simgt 500~\\rm{AU}$, which agrees with the ${}^{13}{\\rm{CO}}$ disk masses of HL Tauri and L1551-IRS5 cited above. Finally, we note that the age of the oldest part of the molecular outflow from L1551-IRS5 is estimated to be $\\sim 10^5~\\rm{yr}$ (Bachiller, Tafalla, \\& Cernicharo 1994). This age is consistent with the time $\\delt$ needed for the central mass in our typical model to become $\\simgt 0.3~\\msol$ (see Fig. $2e$).  Another observational consequence of our model is the magnetic field structure in the core after PMF. Figures $6b$ and $6d$ show that $\\BrZ \\approx \\Beq$ inside the core for the radius range $ 2 \\times 10^{-4} \\simlt r/\\rzero \\simlt 4 \\times 10^{-3}$. Hence, there is significant curvature of field lines (though, as discussed in \\S~3.3, there is less bending than there would be if the field had remained frozen into the neutrals), with bending angles $\\theta_B \\approx \\arctan(\\BrZ/\\Beq)$ in the range $20^{\\circ}- 50^{\\circ}$ for $180~ {\\rm AU} \\simlt r \\simlt 3.5 \\times 10^3~{\\rm{AU}}$. \\footnote{In agreement with the results of CM94 and CM95, we find that the magnetic tension force is generally not negligible in comparison with the magnetic pressure force at any radius (in either the core or the envelope), both before and after PMF.} This type of field geometry may be described roughly as having an hourglass shape. In fact, sub-mm polarimetry of the cores of W3 IRS5 (Greaves, Murray, \\& Holland 1994), Mon R2 (Greaves, Holland, \\& Murray 1995), and OMC-1 (Schleuning 1998) find field geometries suggestive of an hourglass shape on sub-parsec scales. In contrast, field lines remain essentially straight and parallel in the magnetically supported envelope ($r > 0.1~\\rm{pc}$), even after PMF. This agrees with polarimetric observations of molecular clouds that indicate well-ordered fields on these scales (e.g., Hildebrand, Dragovan, \\& Novak 1984; Hildebrand 1989, 1996; Novak et al. 1989; Kane et al. 1993; Hildebrand et al. 1995).  A unique prediction of our model is the large ion--neutral drift speed that occurs during the post-PMF epoch. As shown in Figure $6f$, effective ambipolar diffusion following PMF yields $\\vd \\approx |\\vn| \\gg C$ in the inner regions of the core. In our typical model we find $\\vd \\simgt 1~\\rm{km}~{\\rm s}^{-1}$ for $\\delt \\simgt 2 \\times 10^4~\\rm{yr}$ on scales $r \\simlt 2 \\times 10^{-4} \\rzero \\simeq 180~\\rm{AU}$. Large drift speeds between neutrals and ions (such as $\\HCOp$, $\\rm{HCN}^+$, $\\rm{DCO}^+$, to name but a few) on these scales are therefore expected in our model. Detection of such drift speeds (through high-resolution observations of HCN or $\\HCOp$, say) could be used to observationally confirm our model results and to distinguish them from those of nonmagnetic collapse calculations (e.g., Shu 1977, Hunter 1977, Foster \\& Chevalier 1993) or of magnetic collapse models that do not account for the effect of ambipolar diffusion (e.g., Tomisaka 1996; Li \\& Shu 1997). \\subsection{Features of the Hydromagnetic Shock} The properties of hydromagnetic shocks in partially ionized gases have been developed extensively by many other authors (e.g., Mullan 1971; Draine 1980; Chernoff 1987; Roberge \\& Draine 1990; Draine \\& McKee 1993; Smith \\& Mac Low 1997). Because the ion {\\Alf} speed $v_{\\rm{A,i}} = \\Beq/(4 \\pi \\mi \\nni)^{1/2} = (\\mn/\\mi)^{1/2} (\\nn/\\nni)^{1/2} \\van$ is much larger than $\\van$, $|\\vn|$, and $|\\vi|$ in our model, we expect the outward-propagating shock that develops after PMF to have a magnetic precursor. This is indeed what we find in our simulation: the jump in the ion speed $\\vi$ and the magnetic field strength $\\Beq$ typically occurs at a distance of 1 to 3 computational mesh spacings further away from the symmetry axis than the jump in the neutral infall speed $\\vn$ and the column density $\\sign$. The displacement between the locations of the head of the disturbance in the neutral fluid and in the plasma and magnetic field decreases at later times. Examination of Figures $5b$, $5c$, and $6e$ reveals that, in the reference frame of the shock, the preshock infall speeds are supersonic, and the postshock speeds are also supersonic or just slightly subsonic. Similarly, the ion infall speeds (in the frame of the shock) are much less than the ion {\\Alf} speed. Therefore the shock we observe in our model is probably best classified as being C-type. \\footnote{A C-type shock is characterized by neutral velocities that (in the shock frame) remain supersonic throughout. Hence the shock in our simulation cannot be strictly of this type when the downstream neutral speed is subsonic. In that case a viscous (J-type) subshock may form (Draine \\& McKee 1993), although, as noted by Li \\& McKee (1996), turbulent diffusivity behind a real shock could plausibly keep the postshock flow supersonic and thereby obviate the need for such a subshock.} Making the approximation that in the vicinity of the shock the predominant magnetic stress is that due to the magnetic pressure gradient, the ion force equation becomes \\begin{equation} \\label{ionforceq} \\frac{\\sign}{\\tni} \\vd = - \\frac{Z}{4 \\pi} \\frac{\\partial \\Beq^2}{\\partial r} \\end{equation} (see eqs. [28c] and [51] in CM93, which contain additional terms, involving in particular the magnetic tension force, that could be used to refine the following simple estimate). This yields an approximate shock width \\leteq \\begin{eqnarray} \\label{shkapproxa} \\delshk &\\approx& \\frac{\\Bequ^2}{4 \\pi \\sign} \\frac{\\tni Z}{\\vd} \\left[\\left(\\frac{\\Beqd}{\\Bequ}\\right)^2 -1 \\right] = C \\tni \\frac{\\left(\\vanu/C \\right)^2}{2 \\left(\\vd/C\\right)} \\left[\\left(\\frac{\\Beqd}{\\Bequ}\\right)^2 -1 \\right] \\\\ \\label{shkapproxb} &=& 7.8 \\times 10^{13} \\frac{\\left(T/10~{\\rm K}\\right)^{1/2}\\left(\\vanu/C\\right)^2}{\\left(\\mn/2.33~{\\rm{a.m.u.}}\\right)^{1/2}\\left(\\nni/0.1~\\cc\\right)\\left(\\vd/C\\right)} \\left[1 + 0.067 \\frac{\\left(\\mHII/2~{\\rm{a.m.u.}}\\right)}{\\left(\\mi/30~{\\rm{a.m.u.}}\\right)}\\right] \\nonumber \\\\ &&\\hspace{5em}\\times \\left[\\left(\\frac{\\Beqd}{\\Bequ}\\right)^2-1\\right]~\\rm{cm}\\ , \\end{eqnarray} \\beq where $\\Bequ$ and $\\Beqd$ are the values of $\\Beq$ upstream and downstream of the shock, and $\\vanu$ is the upstream {\\Alf} speed. In deriving the last equality of equation (\\ref{shkapproxa}) we have used the relation $\\sign= 2 \\rhon Z$; equation (\\ref{tnieq}) and the relation $C= \\left(\\kB T/\\mn\\right)^{1/2}$ have been used in deriving equation (\\ref{shkapproxb}). At the time $\\delt_6$ the shock front is located at $r \\simeq 3.9 \\times 10^{-3} \\rzero = 5.2 \\times 10^{16}~{\\rm cm}$, and, in the vicinity of the front, $\\nni \\simeq 10^{-2}~\\cc$, $\\vanu \\simeq C$, $\\Beqd/\\Bequ \\simeq 4.6$ (see Fig. $6b$), and $\\vd \\simeq 0.4 C$ (see Fig. $6f$). For these values our rough estimate for the shock width given by equation (\\ref{shkapproxb}) yields $\\delshk \\simeq 4.3 \\times 10^{16}~{\\rm{cm}}$. Examination of Figure $6e$ at the time $\\delt_6$ reveals that the shock has an actual width $\\delshk \\simeq 2.1 \\times 10^{-3} \\rzero = 2.8 \\times 10^{16}~\\rm{cm}$. (Thus $\\delshk/r_{\\rm shk} \\approx 0.5$ at that time, so the thin-shock approximation that underlies the estimate [17] is marginally satisfied.)  As noted in \\S~1, Li \\& McKee (1996) proposed that a hydromagnetic shock would form in a collapsing core as a result of the decoupling of flux from the ion and neutral fluids because of Ohmic dissipation (a process that becomes important in regions of density $\\nn \\gg 10^{11}~\\cc$; see \\S~2.1). They argued that the accumulating flux diffuses outward to regions of lower density, where improved coupling with the matter causes it to present an obstacle to the infalling neutral gas --- thereby giving rise to a hydromagnetic shock. While our numerical results have confirmed Li \\& McKee's basic shock-formation scenario, they have revealed that ambipolar diffusion, which mediates the shock, can also, following PMF, halt the inward advection of magnetic flux on scales ($r \\simgt 5~\\rm{AU}$) where Ohmic dissipation is not important. In other words, our results have shown that, during the post-PMF epoch, the field--matter decoupling that drives the hydromagnetic shock is due to ambipolar diffusion alone and does not depend on the effect of Ohmic dissipation at $r < \\rinner$.  Because of the similarity in the basic shock-formation mechanism, it is of interest to compare our detailed simulation results with the predictions of the (simplified and analytic) shock model of Li \\& McKee (1996). From their requirement that the magnetic pressure of the shock balance the ram pressure of the neutrals, which were assumed to be freely-falling into the shock, Li \\& McKee derived relations for the shocked magnetic field strength and the shock location (see their eqs. [7] and [8]) in terms of the accretion rate $\\mdot$, the flux-to-mass ratio in units of the critical value for collapse (dubbed $\\epsilon$ in their paper), a parameter related to the logarithmic gradient of the magnetic field (dubbed $\\chi$), the ratio $Z/r$ of the local gas scale height and the radius (dubbed $h$), and the protostellar mass (denoted by $m_{\\ast}$).  At the time $\\delt_6$ we have at the location of our shock $\\dot{M} \\approx 9~\\msol~{\\rm{Myr}}^{-1}$, $\\epsilon \\approx 0.9$, $\\chi \\approx 2$, $h \\approx 0.3$, and $m_{\\ast}=\\mcent(\\delt_6) \\approx 1~\\msol$. Inserting these values into their equations (7) and (8) yields a shocked magnetic field strength $\\sim 630~\\mu{\\rm G}$ and a shock radius $\\sim 2.1 \\times 10^3 ~{\\rm{AU}}$; by comparison, in our model the shocked magnetic field strength at that time is $\\Beqd \\approx 330~\\mu\\rm{G}$ and the shock radius is $r_{\\rm shk} \\approx 3.5 \\times 10^3~\\rm{AU}$ (see Figs. $6b$ and $6e$). The main reason why the analytic expression overestimates the numerically calculated magnetic field strength is that, contrary to the assumption of Li \\& McKee, the preshock neutrals are not in free fall but, rather, are strongly decelerated by magnetic forces (in our simulation we find that the preshock acceleration of the neutrals is reduced to $\\simeq 0.25 \\gr$). The corresponding reduction in the preshock ram pressure leads to a lower value of the postshock field amplitude, with a further reduction in the calculated field strength brought about by the contribution of magnetic tension (ignored in the analytic estimate) to the total magnetic force. Since the analytic estimate of $r_{\\rm shk}$ is based on relating the postshock field strength to the total magnetic flux inside the shock, the overestimate of the field strength naturally results in an underestimate of the shock radius. Despite these discrepancies, Li \\& McKee's analytic representation of the shock parameters provides a decent approximation to the results of our numerical calculation.  Comparison of the column density $\\sign$ (see Fig. $6c$) and neutral infall speed $\\vn$ (see Fig. $6e$) behind the shock shows qualitative agreement with Figures $2a$ and $2d$ of Li \\& McKee (1996), including the free-fall behavior near the central protostar. In the pre-shock region, however, our infall speeds are smaller than theirs because of their assumption of free fall upstream of the shock: Li \\& McKee typically overestimate $\\vnu$ by a factor $\\sim 2$. As a result, the Mach number of the shock relative to the upstream gas (and thus the shock strength) is greater in their model than in ours (see Fig. $5c$) by a similar factor. We have not compared our results for the magnetic field structure behind the shock with those of Li \\& McKee on account of the fact that their system of MHD equations was not closed (it did not include the magnetic induction equation), so that they were unable to calculate the magnetic field with any accuracy (see their Fig. $2b$). \\subsection{Stability of the Core Against Magnetic Interchange} Our simulation has revealed that rapid ambipolar diffusion occurs behind the outward-propagating HMD. The effect that this has on the mass in the flux tubes downstream of the HMD can be seen in Figure $8a$, which shows the local mass-to-flux ratio $d M/d \\phiB = \\sign/\\Beq$ (normalized to the critical value for collapse) as a function of $r/\\rzero$ for the same seven times $\\delt_j$ as in Figure 6. For $\\delt > \\delt_1$ a local minimum in $\\sign/\\Beq$ appears after the passage of the HMD. Hence, there is a region behind the HMD for which $d (\\sign/\\Beq)/dr > 0$. This is a necessary condition for the onset of an interchange instability (e.g., Spruit \\& Taam 1990; Lubow \\& Spruit 1995; Spruit, Stehle, \\& Papaloizou 1995). Li \\& McKee (1996) speculated that such a situation could arise in the wake of a hydromagnetic shock in a collapsing core and suggested that it would act as source of turbulence in the postshock region of the inflow. \\footnote{Li \\& McKee (1996) also noted that the shock may be unstable to the Wardle instability, which involves ions collecting in magnetic field ``valleys'' (Wardle 1990). However, the shock will be immune to this instability if the ion density is determined by the local chemical reaction balance (as assumed in our calculation) rather than by the divergence of the ion mass flux.}  Blaes \\& Balbus (1994) found that the magnetic shearing instability in differentially rotating disks could be stabilized if the disk is weakly ionized. This will also be the case for the interchange instability in a weakly ionized disk: instability is possible only if the growth rate $\\gaminst$ and the neutral--ion collision time $\\tni$ satisfy the condition $\\gaminst \\tni < 1$. This condition reflects the fact that there has to be sufficient collisional coupling between the ion and neutral fluids for a magnetic interchange instability to grow in the neutrals; otherwise the instability is damped. Calculation of $\\gaminst$ in our model is hampered by the fact that previous studies of interchange instability have been carried out only for disks that are in hydrostatic equilibrium, with exact balance between magnetic and gravitational forces. We can apply the results of these studies to our model only if the region behind the shock where $d (\\sign/\\Beq)/dr > 0$ is effectively in quasi equilibrium, with approximate balance between gravitational and magnetic forces, and with infall speeds $|\\vn| \\ll (r |\\gr|)^{1/2}$ ($\\simeq$ the free-fall speed). In our model, the magnitude of the acceleration of the neutrals in this region of the core does not become $\\simlt 0.1 | \\gr |$ until times $\\sim \\delt_6$; hence, approximate equilibrium between magnetic and gravitational forces is valid only at these later times. Spruit \\& Taam (1990) found that the growth rate for the most unstable linear interchange modes is \\begin{equation} \\label{growtheq} \\gaminst = \\left( \\frac{\\Beq \\BrZ}{2 \\pi \\sign} \\frac{d}{dr} \\ln \\frac{\\sign}{\\Beq} \\right)^{1/2} . \\end{equation} The product $\\gaminst \\tni$ for the region of the core susceptible to interchange instability is shown in Figure $8b$ as a function of $r/\\rzero$ at the time $\\delt_6$. We also plot (Fig. $8c$) the product $\\gaminst \\tau_{\\rm{kin}}$, where $\\tau_{\\rm{kin}} \\equiv r/|\\vn|$ is the kinematical timescale, as a function of $r/\\rzero$ for the same region of the core and time. If this product is $ < 1$, the unstable modes will be ``swept up'' by the infalling gas before they have time to grow. It is evident from these figures that $\\gaminst \\tni < 1$ and $\\gaminst \\tau_{\\rm{kin}} > 1$ for the region of the core susceptible to magnetic interchange. This means that there is sufficient collisional coupling between the ions and neutrals, and that the instability will grow before being swept along with the neutrals. Hence, this region of the core may be interchange unstable. An instability of this type would enhance the tansfer of gas with a high mass-to-flux ratio to the center (e.g., Spruit \\& Taam 1990), and, as noted by Li \\& McKee (1996), might also lead to the development of turbulence that could increase the field diffusivity in the postshock gas. However, the onset and development of this instability can only be studied by means of a fully 3-D simulation. \\subsection{Implications to the Magnetic Flux Problem} The magnetic flux problem in star formation has to do with the fact that the magnetic flux of a $1\\msol$ blob of matter in the diffuse interstellar medium is typically several orders of magnitude greater than the flux of a $1\\msol$ protostar. Such a blob of matter would therefore have to get rid of most of its flux before becoming a star. Ambipolar diffusion has long been suggested as a means by which the magnetic flux problem could be resolved (e.g., Mestel \\& Spitzer 1956; Mouschovias 1978; Paleologou \\& Mouschovias 1983; Nakano 1984; Mouschovias, Paleologou, \\& Fiedler 1985). In general, these earlier studies focused primarily on the role of ambipolar diffusion and the magnetic flux problem for the pre-PMF phase of protostellar evolution. \\footnote{On the basis of a consideration of the timescales for ambipolar diffusion and Ohmic dissipation at high densities, Nakano \\& Umebayashi (1986b) suggested that significant flux loss could only take place (primarily by Ohmic dissipation, according to their estimates) during the dynamical phase of core collapse. Lizano \\& Shu (1989) similarly concluded that the resolution of the protostellar magnetic flux problem must occur during the dynamical stage of core evolution: using the quasi-static approximation (valid for $\\nnc \\simlt \\rm{a~few} \\times 10^4~\\cc$; Fiedler \\& Mouschovias 1993; CM94; Basu \\& Mouschovias 1994) to calculate the contraction of a slightly subcritical molecular cloud, they found that only a small amount of flux is lost by ambipolar diffusion from the central flux tubes before runaway collapse is initiated.} While ambipolar diffusion does indeed reduce the flux-to-mass ratio during that phase, the flux contained within a $1\\msol$ region of a molecular cloud core is still much larger at the time of PMF than typical protostellar fluxes. Specifically, CM94 and CM95 found that the central $1\\msol$ flux tube within their cores had a total magnetic flux $\\sim 10^{30}~\\rm{Mx}$ (consistent with our results in \\S~3.2) during the pre-PMF dynamical collapse phase of the typical model. This value represents a reduction by a factor $\\sim 5.6$ of the flux associated with that mass before the onset of ambipolar diffusion. Nevertheless, it greatly exceeds the plausible upper limit ($\\sim 6 \\times 10^{26}~\\rm{Mx}$) on the flux of a solar-mass protostar (estimated assuming an average surface field of 10 kG and a stellar radius of $10^{11}~{\\rm cm}$; see Li \\& McKee 1996).  We have shown in this paper that the rate of ambipolar diffusion is strongly increased during the post-PMF epoch of star formation. It is therefore of interest to examine the implications of our simulation results to the magnetic flux problem. As discussed in \\S~2.2, in our calculations we only consider the core regions at radii $r \\simgt \\rinner$ ($\\simeq 7.3~\\rm{AU}$ for our typical model), where ambipolar diffusion is the dominant mechanism of flux loss. Initially, the magnetic flux contained within $\\rinner$ is $\\phicent(\\delt=0) = 6.7 \\times 10^{26}~\\rm{Mx}$. As shown in Figure $3a$, the central flux increases before the onset of rapid ambipolar diffusion. This continues to the time $\\delt \\approx 10^3~\\rm{yr}$. For $\\delt > 10^3~\\rm{yr}$, ambipolar diffusion prevents further advection of flux from $r > \\rinner$ into the central sink, and $\\phicent$ changes very little after this time. By the time $\\delt \\approx 10^5~\\rm{yr}$, when $\\mcent \\approx 1\\msol$, $\\phicent \\approx 5 \\times 10^{27}~\\rm{Mx}$. This represents a decrease of over two orders of magnitude relative to the flux associated with this mass at the time of PMF. While this value is still about an order of magnitude higher than our adopted upper limit on the protostelar flux, the discrepancy is now much lower than previous estimates of ambipolar diffusion have indicated. The important conclusion from our work is thus that {\\em ambipolar diffusion in contracting molecular cloud cores can in principle contribute significantly to the resolution of the magnetic flux problem} by reducing the magnetic flux brought into a solar-mass protostar by a factor $\\simgt 10^3$. The new, and somewhat surprising, result is that {\\em most of this reduction occurs after PMF.}  The remainder of the protostellar magnetic flux could possibly be extracted from the infalling mass through Ohmic dissipation within $\\rinner$, although refreezing of the magnetic field into the matter, brought about by collisional reionization at densities $\\nn \\simgt 10^{14}~\\cc$ (e.g., Pneuman \\& Mitchell 1965; Nakano \\& Umebayashi 1986b; Li \\& McKee 1996), as well as anomalous diffusivity (e.g., Norman \\& Heyvaerts 1985) operating in the reionized gas, could complicate the issue. Another complicating factor is the strong likelihood that much of the mass and flux carried into the protostar pass through a rotationally supported, circumstellar accretion disk of size $\\gg \\rinner$ (e.g., Lubow, Papaloizou, \\& Pringle 1994; Reyes-Ruiz \\& Stepinski 1996; Li 1996). It is also conceivable that magnetic flux is brought to the vicinity of the protostar but excluded from its interior by turbulent diffusivity associated with convection. Since the region within $\\rinner$ was excluded from our calculation, we do not pursue this topic any further in this paper."
    },
    "9804/hep-ex9804007_arXiv.txt": {
        "abstract": "With an effective telescope area of order $10^4$~m$^2$, a threshold of $\\sim$50~GeV and a pointing accuracy of 2.5~degrees, the AMANDA detector represents the first of a new generation of high energy neutrino telescopes, reaching a scale envisaged over 25 years ago. We describe its performance, focussing on the capability to detect halo dark matter particles via their annihilation into neutrinos. ",
        "introduction": "\\unskip  High energy neutrino telescopes are multi-purpose instruments; their science mission covers particle physics, astronomy and astrophysics, cosmology and cosmic ray physics. Their deployment creates new opportunities for glaciology and oceanography, possibly geology of the earth's core\\cite{pr}. Astronomy with neutrinos does have definite advantages. They can reach us, essentially without attenuation in flux, from the largest red-shifts. The sky is, in contrast, partially opaque to high energy photons and protons because of energy-loss suffered in interactions with infrared light, CMBR photons and radio waves\\cite{cronin}. They do not reach us from distances much larger than tens of megaparsecs once their energy exceeds thresholds of 10~TeV for photons and $5\\times 10^7$~TeV for protons. (Below this energy charged protons do not point back to their sources.) The drawback is that neutrinos are difficult to detect: the small interaction cross sections that enable them to travel without attenuation over a Hubble radius, are also the reason why kilometer-scale detectors are required in order to capture them in sufficient numbers to do astronomy\\cite{halzenkm}. Some opportunities may, however, be unique. If, for instance, the sources of the highest energy cosmic rays are beyond $10^2$~Mpc, conventional astronomy is unlikely to discover them.  Some science missions do not require a detector of kilometer size. The best opportunities to search for halo dark matter are, in fact, associated with the present instrument which, while smaller in telescope area than the planned extension of AMANDA to ICE3 (ICECUBE), has a lower threshold.   At this meeting, the capability of neutrino telescopes to discover the particles that constitute the dominant, cold component of the dark matter is of special interest. The existence of the weakly interacting massive particles (WIMPs) is inferred from observation of their annihilation products. Cold dark matter particles annihilate into neutrinos; {\\it massive} ones will annihilate into {\\it high-energy} neutrinos which can be detected in high-energy neutrino telescopes. This so-called indirect detection is greatly facilitated by the fact that the earth and the sun represent dense, nearby sources of accumulated cold dark matter particles. Galactic WIMPs, scattering off nuclei in the sun, lose energy. They may fall below escape velocity and be gravitationally trapped. Trapped WIMPs eventually come to equilibrium and accumulate near the center of the sun. While the WIMP density builds up, their annihilation rate into lighter particles increases until equilibrium is achieved where the annihilation rate equals half of the capture rate. The sun has thus become a reservoir of WIMPs which we expect to annihilate mostly into heavy quarks and, for the heavier WIMPs, into weak bosons. The leptonic decays of the heavy quark and weak boson annihilation products turn the sun and earth into nearby sources of high-energy neutrinos with energies in the GeV to TeV range. Figure~1 displays the neutrino flux from the center of the earth calculated in the context of supersymmetric dark matter theories\\cite{scopel}. The direct capture rate of the WIMPs in germanium detectors is shown for comparison. Contours indicate the parameter space favored by grand unified theories. Most of this parameter space can be covered by improving the capabilities of existing detectors by 2 orders of magnitude. Existing neutrino detectors have already excluded fluxes of neutrinos from the earth's center of order 1~event per $1000 \\rm~m^2$ per year. The best limits have been obtained by the Baksan experiment\\cite{suvorova}. They are already excluding relevant parameter space of supersymmetric models. We will show that, with data already on tape, the AMANDA detector will have an unmatched discovery reach for WIMP masses in excess of 100~GeV.  \\begin{figure}[t] \\centering \\hspace{0in}\\epsfxsize=4.5in\\epsffile{fig1.eps}  \\caption{Direct and indirect detection rates (for neutrinos from the center of the earth in the figure shown) of cold dark matter particles predicted by supersymmetric theory. Grand unified theories favor the parameter space indicated. Part of it is already excluded by present experiments as indicated by the horizontal and vertical lines.} \\end{figure}  The potential of neutrino telescopes as dark matter detectors has been documented in detail\\cite{kamionkowski}. With a sensitivity which increases with the WIMP mass, they are complementary to direct, cryogenic detectors. They can detect WIMPS beyond the kinematic limits of the LHC: about 500~GeV for neutralinos. A striking way to illustrate their potential is to use the possible detection\\cite{belli} in the DAMA NaI detector in the Gran Sasso tunnel as an example. If their seasonal variation is indeed evidence for WIMPS, observation of a signal in an exposure of 4500 kg\\,days requires a WIMP-nucleon cross section of $10^{-42}{\\sim}10^{-41}$~cm$^2$ for a WIMP mass of $50{\\sim} 150$~GeV. This information is sufficient to calculate their trapping and annihilation rate in the sun and earth. Both will be a source of, on average, 100 neutrinos per year of WIMP origin in the existing AMANDA detector with an effective area of $10^4$~m$^2$. The exact rate varies with the mass of the WIMP. ",
        "conclusions": ""
    },
    "9804/astro-ph9804107_arXiv.txt": {
        "abstract": "We have computed line profiles from self-gravitating toroids around black holes. The specific angular momentum of the toroids is assumed to be constant in space. The images of the toroids show peculiar feature in the rear side of the black holes. Concerning the line profiles, the red wing extends to the very small frequency region because the location of the inner edge is rather near the event horizon of the black hole and consequently the velocity of the inner edge of toroids can be faster than that of the Kepler disks. ",
        "introduction": "Recently, strong evidence for existence of supermassive black holes in the central regions of galaxies has been found from observations of electromagnetic waves with various wavelengths ranging from the radio waves to X-rays (see e.g. Miyoshi et al. 1995 for the radio observation; Ford et al. 1994, Harms et al. 1994 for the optical observation; Tanaka et al. 1995 for the X-ray observation).  Radio and optical observations have shown that there exist very rapidly rotating gaseous disks in the central regions of galaxies. Since the high velocity of a disk implies the existence of a large amount of mass within a region of a very small size, it is widely considered that there are supermassive black holes with masses of $10^7 \\sim 10^9 M_{\\sun}$ at the centers of galaxies. However, such observations do not reveal the nature of black holes because the size of the observed region is still too large to get detailed information about the black holes.  Contrary to these optical and radio analyses, recent X-ray observations have brought us important information about gravitational fields very near the massive objects or black holes. By using the ASCA satellite, broad iron emission lines have been detected in active galaxies (see e.g. Fabian et al. 1994; Mushotzky et al. 1995). In particular, Tanaka et al.~(1995) observed the Seyfert 1 galaxy, MCG--6--30--15, and discovered a very broad and skewed iron emission line. The broadness and skewness of the line profile can be explained only by assuming that the inner edge of the accretion disk is located very near the event horizon of the central black hole (Tanaka et al. 1995; Fabian et al. 1995). It implies that X-ray observations can be used to understand the very vicinity of the black holes. It is important to know the gravitational field near the event horizon because it is the key feature to distinguish a Schwarzschild black hole from a Kerr black hole.  Consequently we have reached a stage to investigate very strong gravitational fields of black holes and to be able to test the validity of general relativity.  However, since the emission line profiles depend not only on the gravitational fields but also on the structures of the accretion disks, it is not easy to determine the type of the black hole, i.e. whether the central black hole is of a Schwarzschild type or of a Kerr type. In fact, concerning the nature of the black hole at the center of MCG--6--30--15, various discussions have not settled down to a unique interpretation yet (see e.g. Tanaka et al. 1995; Iwasawa et al. 1996; Dabrowski et al. 1997; Bromley et al. 1997; Reynolds \\& Begelman 1997; Bromley et al. 1998). Therefore, in order to get a consistent picture of black hole -- accretion disk systems, we have to pursue much more investigations theoretically as well as observationally.  Concerning the theory about the spectra of accretion disks, Cunningham~(1975) was the first to formulate the problem for Kerr black holes and obtained theoretical spectra from accretion disks around black holes (see also Cunningham 1976).  Gerbal \\& Pelat~(1981) investigated lines from a ring around a black hole and found double-peaked asymmetric profiles. From the end of 80's, observations of lines in the X-ray spectrum stimulated many authors to study line profiles of accretion disks as well as disk structures (e.g. Nandra et al. 1989; Fabian et al. 1989; Kunieda et al. 1990; Kojima 1991; Laor 1991; Chen \\& Halpern 1989, 1990).  However, in almost all theoretical studies mentioned above, investigations of the emission line profiles have been done by assuming that disks are geometrically thin and that only direct photons are observed.  Some authors have studied the effect of multiple images (Luminet 1978; Bao, Hadrava \\& {\\O}stgaard 1994) and that of self-eclipse due to toroidal configurations (Bao \\& Stuchlik 1992; Kojima \\& Fukue 1992).  For toroidal configurations the rotation law of the toroid is not always that of the Kepler rotation because of the presence of the pressure within the toroid. In fact, Kojima \\& Fukue~(1992) employed a variety of rotation laws. However, their analysis was done in the framework of Newtonian gravity. Therefore, quantitative treatment of non-Keplerian toroids in general relativity has not been carried out yet.  Furthermore, in some situations, self-gravity of disks or toroids plays an important role for the structures of disks or toroids. In particular, massive neutron toroids around neutron stars or black holes have been proposed as possible sources of $\\gamma$-ray bursts (e.g. Paczy\\'nski 1991; Narayan, Paczy\\'nski \\& Piran 1992; Jaroszy\\'nski 1993; Witt et al. 1994).  Although there might be little chance to observe line profiles from such systems even if exist,  it would be interesting to study the effect of self-gravity of toroids on the line profiles.  Such self-gravitating disks were investigated by Karas, Lanza \\& Vokrouhlicky~(1995). However, they have studied very light thin disks whose mass is less than several percent of the mass of the black hole. From the theoretical point of view, it would be interesting to investigate more massive disks or toroids as well as the effect of geometrical thickness. By treating massive toroids, the central objects are no more of Schwarzschild type nor Kerr type black holes because of the gravitational effect of the self-gravitating massive toroids on the black holes (Nishida \\& Eriguchi 1994). Moreover, gravitational effect of massive toroids may bring some differences to line profiles.  Therefore, in this paper, we will consider self-gravitating toroid -- black hole systems and study their effect on emission line profiles. ",
        "conclusions": "In this paper we have constructed toroid -- black hole systems in general relativity. We have computed photon trajectories in the numerically obtained gravitational fields and investigated the images of the toroids as well as the line profiles from the toroids. In these computations, we have assumed that the observer is located very near the black hole because the metric in the whole space has not been calculated. Therefore, the quantitative values in this paper will be changed a little if the observer is located at infinity. However, the important thing is not to obtain exact quantitative values but to know characteristic feature which appears only by introducing thickness and self-gravity numerically exactly. In this sense, our results serve as representative ones for self-gravitating toroids around black holes.  As discussed in Introduction, the observation of MCG--6--30--15 has given us time varying line profiles whose interpretation has not been clarified yet (Tanaka et al. 1995; Iwasawa et al. 1996; Dabrowski et al. 1997). Some authors have proposed models for the system which would explain the observations (Dabrowski et al. 1997; Bromley et al. 1997; Bromley et al. 1998). However, since there are many parameters about the disk structures and the X-ray sources, it is very difficult to obtain a unique solution for the system.  This can be also seen from our result. For our Model B with the emissivity indices $s = -1$ and $s = -4$, the line profiles are shown in Fig.~\\ref{comparison} together with the observational data. As seen from this figure, the tendency of the observational data seems to be roughly explained by the change of the emissivity index of Model B.  Thus, at the present stage, we can only say that since the model cannot be uniquely determined, we have to get more accurate observational data to clarify the spacetime of the central region of galaxies.   \\begin{figure*} \\begin{tabular}{cc} (a)&(b)\\\\ \\epsfile{file=tanaka.eps,width=0.5\\textwidth} &\\epsfile{file=iwasawa.eps,width=0.5\\textwidth}\\\\ \\end{tabular} \\caption{ Same as Fig. 4 but for the model with $\\theta_0 = 30 \\degr$ and emissivity indices (a) $s=-1$ (solid line) and (b) $s=-4$ (solid line). The observational data of the ASCA are plotted by points with error bars for the data of Tanaka et al.~(1995) (a) and that of Iwasawa et al.~(1996) (b). }   \\label{comparison} \\end{figure*}"
    },
    "9804/astro-ph9804331_arXiv.txt": {
        "abstract": "We report the detection of the GRB 971214 counterpart in the near-infrared by means of two images in the K$^{\\prime}$-band taken at Calar Alto only $\\sim$3.5 and $\\sim$5 hours after the gamma-ray event. We detect the transient at K$^{\\prime}$=$18.03\\pm0.18$ and K$^{\\prime}$=$18.00\\pm0.22$ respectively. Our data seem to indicate the existence of a plateau with duration $ 1.5 \\leq T \\le 6.7 $ hours (between 3.5 and 10.2 hours after the high-energy event). Moreover the power-law decline should be steeper than the one given by the index $\\alpha_{\\rm K^{\\prime}}=0.45$. There is also a change in the slope of the broad-band spectrum at some wavelength between the J and K$^{\\prime}$ bands (possibly around the H-band). ",
        "introduction": "After 31 years since the discovery of gamma-ray bursts (GRBs), the origin of such brief gamma-ray flashes remains unknown. The observed isotropy of GRBs in the sky, could only be explained by theoretical models where GRBs originate either in a extended halo around the Galaxy or arise from sources at cosmological distances. Before the launch of the BeppoSAX satellite, the poor localization capability of the GRB detectors made the searches at other wavelengths unfruitful.  The breakthrough took place in 1997 when the first X-ray afterglows were observed by the BeppoSAX, RXTE, ROSAT and ASCA satellites (Costa et al. 1997, Heise et al.  1997a, Marshall et al. 1997, Greiner et al. 1998, Murakami 1998).  They were able to localize the fading X-ray emission that followed the more energetic gamma-ray photons once the GRB event had ended. This emission (the afterglow) extends to longer wavelengths, and the good accuracy in the position determination by BeppoSAX (typically $1^{\\prime}$ radius error boxes) has led to the discovery of the first optical counterparts for GRB 970228 (van Paradijs et al. 1997, Guarnieri et al. 1997), and GRB 970508 (Bond 1997, Djorgovski et al.  1997, Castro-Tirado et al.  1998), greatly improving our understanding of these puzzling sources. The measurement of the redshift for the GRB 970508 optical counterpart (Metzger et al.  1997) has established that one GRB, maybe all, lie at cosmological distances.  GRB 971214 is the third GRB with a known optical counterpart. It was detected by the BeppoSAX Gamma-ray Burst Monitor (GBM, Frontera et al. 1997) on Dec 14.97 1997, as a 25 s long-structured gamma-ray burst. Simultaneous to the detection of the GBM, the Wide Field Cameras (WFC, Jager et al.  1997) on board BeppoSAX provided an accurate position (a $3^{\\prime}.9$ radius error box at a 3$\\sigma$ confidence level, Heise et al. 1997b) that allowed deep optical, infrared and radio observations. The position was also consistent with the one given by the all-sky monitor on RXTE (Bradt et al. 1993) and by the BATSE/Ulysses annulus (Kippen et al. 1997). When BeppoSAX pointed its Narrow-Field Instruments (NFI) to the GRB position, on Dec 15.25 ($\\sim$ 6.5 hours after the burst), a previously unknown variable X-ray source was found inside the WFC error box (Antonelli et al.  1997) which was identified as the X-ray afterglow of GRB 971214. Soon after, Halpern et al.  (1997) reported the presence of a fading object inside the WFC GRB error box, based on two I-band images separated 24 hours. The object was afterwards confirmed as the counterpart of GRB 971214 by means of additional observations at other wavelengths: R-band (Castander et al.  1997, Diercks et al.  1998), I-band (Rhoads 1997) and J-band (Tanvir et al.  1997). No detections in the K-band were reported in the literature, although observations performed on Dec 15.54 imposed an upper limit of K $>$ 18.5 (Garcia et al.  1997).  As it will be explained later, this upper limit will be used to constraint the power-law index $\\alpha_{\\rm K^{\\prime}}$ and the position of the possible maximum of the light curve.   We report here the detection of the GRB 971214 counterpart in the near infrared (IR) by means of two K$^{\\prime}$-band images taken at Calar Alto on Dec 15.12 and 15.18 (mean observing time, only $\\sim$3.5 and $\\sim$5 hours after the gamma-ray event). The second image is almost simultaneous to the beginning of the observations performed by the BeppoSAX narrow-field instruments.  We discuss whether our observations are in agreement with the extrapolation of the power-law seen at other bands in later epochs. ",
        "conclusions": "We have detected the GRB 971214 near-IR counterpart $\\sim$3.5 hours and $\\sim$5 hours after the gamma-ray event which enables to conclude that:  i) a magnitude difference $\\Delta$K${^{\\prime}}=-0.03\\pm 0.28$ is derived from our measurements, whereas $\\Delta$K${^{\\prime}}=0.464$ would be expected assuming a power-law decay with index $\\alpha_{\\rm K^{\\prime}}=1.2$ (similar to the one observed at optical wavelengths). This implies a deviation of $1.7\\sigma$. If the assumed power-law index $\\alpha_{\\rm K^{\\prime}}$ were 1.4, then the rejection level would be $2.0\\sigma$.  Thus, our measurements suggest a rising or a flat light curve segment with a duration $ 1.5 \\leq T \\le 6.7 $ hours (between 3.5 and 10.2 hours after the burst).  This conclusion must be taken with care since the above-mentioned rejection levels are not stringent enough to assure the result with total confidence;   ii) the power-law decline in the near-IR should be steeper than the one given by $\\alpha_{\\rm K^{\\prime}}=0.45$;  iii) for the observations carried out on Dec 15.44-15.51, there is a change in the slope of the measured energy distribution at some wavelength between the J and K$^{\\prime}$ bands (possibly around H)."
    },
    "9804/astro-ph9804041_arXiv.txt": {
        "abstract": "The measurements of CMB anisotropy have opened up a window for probing the global topology of the universe on length scales comparable to and beyond the Hubble radius. For compact topologies, the two main effects on the CMB are: (1) the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the manifold and (2) an infrared cutoff in the power spectrum of perturbations imposed by the finite spatial extent.  We present a completely general scheme using the {\\em regularized method of images} for calculating CMB anisotropy in models with nontrivial topology, and apply it to the computationally challenging compact hyperbolic topologies. This new technique eliminates the need for the difficult task of spatial eigenmode decomposition on these spaces. We estimate a Bayesian probability for a selection of models by confronting the theoretical pixel-pixel temperature correlation function with the {\\sc cobe--dmr} data. Our results demonstrate that strong constraints on compactness arise: if the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the observations. If the universe is of comparable size, the likelihood function is very dependent upon orientation of the manifold {\\it wrt} the sky. While most orientations may be strongly ruled out, it sometimes happens that for a specific orientation the predicted correlation patterns are preferred over the conventional infinite models. ",
        "introduction": "The remarkable degree of isotropy of the cosmic microwave background (CMB) points to homogeneous and isotropic Freidmann-Robertson-Walker (FRW) models for the universe. The underlying Einstein's equations of gravitation are purely local, completely unaffected by the global topological structure of space-time.  In fact, in the absence of spatially inhomogeneous perturbations, a FRW model predicts an isotropic CMB regardless of the global topology.  The observed large scale structure in the universe implies spatially inhomogeneous primordial perturbations exist which gave rise to the observed anisotropy of the CMB. The global topology of the universe does affect the local observable properties of the CMB anisotropy. In compact universe models, the finite spatial size usually precludes the existence of primordial fluctuations with wavelengths above a characteristic scale related to the size of the universe.  As a result, the power in the CMB anisotropy is suppressed on large angular scales. Another consequence is the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the manifold. One can search for such patterns statistically in the COBE maps, and to the extent that they are not there, one can constrain the size of the universe and its topology.  Much recent astrophysical data suggest the cosmological matter density parameter, $\\Omega_0$, is subcritical~\\cite{opencase}. If a (possibly varying) cosmological constant is absent or insufficient to bring the total density to the critical value, this would imply a hyperbolic spatial geometry for the universe (commonly referred to as the `open' universe in the cosmological literature). The topologically trivial (simply connected) hyperbolic 3-space, $\\hm$, is non-compact and has infinite size.  There are numerous theoretical motivations, however, to favor a spatially compact universe~\\cite{motive}.  To reconcile a compact universe with a flat or hyperbolic geometry, consideration of spaces with non trivial topology (non simply connected spaces) is required. A compact cosmological model can be constructed by identifying points on the standard infinite flat or hyperbolic FRW spaces by the action of a suitable discrete subgroup, $\\Gamma$, of the full isometry group, $G$, of the FRW space. The infinite FRW spatial hypersurface is the {\\em universal cover}, tiled by copies of the compact space (most appropriately represented as the {\\em Dirichlet domain} with the observer at its {\\em basepoint}).  Any point ${{\\bf x}}$ of the compact space has an image ${{\\bf x}}_i = \\gamma_i {{\\bf x}}$ in each copy of the Dirichlet domain on the universal cover, where $\\gamma_i \\in \\Gamma$.  The hyperbolic manifold, $\\hm$, can be viewed as a hyperbolic section embedded in four dimensional flat Lorentzian space. The isometry group of $\\hm$ is the group of rotations in the four space -- the proper Lorentz group, $SO(3,1)$. A compact hyperbolic (CH) manifold is then completely described by a discrete subgroup, $\\Gamma$, of the proper Lorentz group, $SO(3,1)$. The Geometry Centre at the University of Minnesota has a large census of CH manifolds and public domain software SnapPea~\\cite{Minn}. We have adapted this software to tile $\\hm$ under a given topology using a set of generators of $\\Gamma$. The tiling routine uses the generator product method and ensures that all distinct tiles within a specified tiling radius are obtained.   A CH manifold, ${\\cal M}$, is characterized by a dimensionless number, ${\\cal V}_{\\!\\cal M}\\equiv V_{\\!\\cal M}/d_c^3$, where $V_{\\!\\cal M}$ is the volume of the space and $d_c$ is the curvature radius \\cite{Thur7984}. There are a countably infinite number of CH manifolds with no upper bound on ${\\cal V}_{\\!\\cal M}$.  The smallest CH manifold discovered so far has ${\\cal V}_{\\!\\cal M} =0.94$~\\cite{smallestCH}. ~\\footnote{ The volume of CH manifolds is bounded from below and the present theoretical lower bound stands at ${\\cal V}_{\\!\\cal M} \\ge 0.167$~\\cite{gab_mey96}. There exist sharper lower bounds within subclasses of CH manifolds under restrictions on topological invariants~\\cite{minvol}. It has been conjectured that the smallest known manifold is in fact the smallest possible~\\cite{smallestCH}.} The Minnesota census lists several thousands of these manifolds with ${\\cal V}_{\\!\\cal M}$ up to $\\sim 7$. In the cosmological context, the physical size of the curvature radius $d_c$ is determined by the density parameter and the Hubble constant $H_0$: $d_c=(c/H_0)/\\sqrt{1-\\Omega_0}$. The physical volume of the CH manifold with a given topology, \\ie a fixed value of $V_{\\!\\cal M}/d_c^3$, is smaller for smaller values of $\\Omega_0$. Two quantities which characterize linear dimensions of the Dirichlet domain are $R_>$ and $R_<$, the radii of circumscribing and inscribing spheres, respectively.  In the standard picture, the CMB that we observe is a Planckian distribution of relic photons which decoupled from matter at a redshift $\\approx 1100$. These photons have freely propagated over a distance $ R_{\\sc ls} \\approx 2 d_c\\, {\\rm arctanh} \\sqrt{1-\\Omega_0} $, comparable to the ``horizon'' size.  For the adiabatic fluctuations we consider here, the dominant contribution to the anisotropy in the CMB temperature measured with wide-angle beams ($\\theta_{\\sc fwhm} \\gta 2^\\circ \\Omega_0^{1/2}$) comes from the cosmological metric perturbations through the Sachs-Wolfe effect.  The adiabatic cosmological metric perturbations can be expressed in terms of a scalar gravitational potential, $\\Phi({\\bf x},\\tau)$.  The dynamical equation for the gravitational potential allows for separation of the spatial and temporal dependence,\\footnote{At the scales appropriate to CMB anisotropies, damping effects on $\\Phi$ can be neglected.}  $\\Phi({\\bf x},\\tau) = F(\\tau) \\Phi({\\bf x},\\tauls)$, where $F(\\tau)$ encodes the time dependence of the metric perturbations and $\\Phi({\\bf x},\\tauls)$ is the field configuration on the three-hypersurface of constant time $\\tau=\\tauls$ when the last scattering of CMB photons took place. We shall study open, $\\Omega_0<1$, models with zero cosmological constant where in the matter dominated phase,~\\cite{mukh92}  \\begin{equation} F(\\tau) = \\frac{5(\\sinh^2\\tau -3 \\tau\\sinh\\tau +4\\cosh\\tau -4)} {(\\cosh\\tau -1)^3}\\,. \\lbl{Feta} \\end{equation} Here and further on we use dimensionless conformal time $\\tau$ expressed in units of the curvature radius. A non-zero cosmological constant can be trivially incorporated in our analysis by using the appropriate solution for $F(\\tau)$.  We write the Sachs-Wolfe formula for the CMB temperature fluctuation, $\\Delta T(\\hat q)$, in a direction $\\hat q$, in the form \\begin{equation} \\fl \\dT(\\hat q) = \\frac{1}{3} \\Phi(\\hat q\\chiH,\\tauls) + 2 \\int_{0}^\\chiH d\\chi f(\\chi) \\Phi(\\hat q\\chi,\\tauls)\\,, \\quad f(\\chi)=\\frac{d}{d\\tau} F(\\tau)\\bigg|_{\\tau=\\chiH-\\chi}\\,, \\lbl{dTSW} \\end{equation} where $\\chi$ is the affine parameter along the photon path from $\\chi=0$ at the observer position to $\\chiH=R_{\\sc ls}/d_c$. The first term is called the {\\em surface} or ``naive'' Sachs-Wolfe effect (NSW). The second term, which is nonzero only if $\\Phi$ varies with time between $\\tauls$ and now, is the {\\em integrated} Sachs-Wolfe effect (ISW). The angular correlation between the CMB temperature fluctuations in two directions in the sky is then given by \\begin{eqnarray} \\fl C(\\hat q,\\hat q^\\prime)\\equiv \\left\\langle\\dT(\\hat q)\\dT(\\hat q^\\prime)\\right\\rangle = \\frac{1}{9} \\langle\\Phi(\\hat q\\chiH,\\tauls)\\Phi(\\hat q^\\prime\\chiH,\\tauls)\\rangle \\nonumber\\\\ \\lo +\\frac{2}{3}\\int_{0}^\\chiH d\\chi~f(\\chi)~\\left[ \\langle\\Phi(\\hat q\\chi,\\tauls)\\Phi(\\hat q^\\prime\\chiH,\\tauls)\\rangle +\\langle\\Phi(\\hat q^\\prime\\chi,\\tauls)\\Phi(\\hat q\\chiH,\\tauls)\\rangle\\right] \\nonumber\\\\ \\lo +4\\int_{0}^\\chiH d\\chix{1} ~f(\\chix{1})~ \\int_{0}^\\chiH d\\chix{2}~f(\\chix{2})~ \\langle\\Phi(\\hat q\\chix{1},\\tauls)\\Phi(\\hat q^\\prime\\chix{2},\\tauls)\\rangle\\,.  \\lbl{cthetaSW} \\end{eqnarray} The main point to be noted is that $C(\\hat q,\\hat q^\\prime)$ depends on the spatial two point correlation function, $\\xi_\\Phi \\equiv \\langle\\Phi({\\bf x},\\tauls)\\Phi({\\bf x^\\prime},\\tauls)\\rangle $ of $\\Phi$ on the three-hypersurface of last scattering. This is due to the fact that the equation of motion for $\\Phi$ allows a separation of spatial and temporal dependence.  Although in this work we restrict our attention to the Sachs-Wolfe effect which dominates when the beam size is large, we should point out that other effects which contribute to the CMB anisotropy at finer resolution can also be approximated in terms of spatial correlation of quantities defined on the hypersurface of last scattering~\\cite{us_inprep}. ",
        "conclusions": ""
    },
    "9804/astro-ph9804088_arXiv.txt": {
        "abstract": "We have imaged the inner square arcminute of the well known lensing and cooling flow cluster A2390 ($z = 0.23$) down to a sensitivity of 65 and 130 $\\mu$Jy at 6.75 and 15 $\\mu$m , respectively. We report the first evidence of an active star-forming region in a cooling flow (at those wavelengths) and strong emission in the mid-IR from lensed galaxies located at $z=0.9$. ",
        "introduction": "The cluster of galaxies A2390 (z=0.228) possesses remarkable properties which makes its study particularly attractive: presence of a ``straight'' giant gravitational arc (z=0.913), numerous arclets, an elongated galaxy distribution (\\cite{Mellier} ; \\cite{Pello}) and a large velocity dispersion (1093 km s$^{-1}$, Carlberg et al., 1996) as well as a high X-ray luminosity ($\\sim1.5 ~10^{45}$ erg s$^{-1}$ in the [0.1--2.4] keV band). A deep HRI ROSAT pointing revealed an elongated X-ray morphology, the existence of a secondary maximum responsible for the observed  gravitational shear in the optical and a strong cooling flow of $\\sim800$ M$_{\\odot}$yr$^{-1}$ (Pierre et al., 1996). All this indicates that A2390, and its underlying gravitational potential, is especially relevant for our understanding of massive cluster formation, which is, in a hierarchical scenario, closely related to the history of galaxy/star formation. This has motivated deep ISOCAM observations of the cluster core  during the guaranteed time programme DEEPXSRC. We present here the observations and  results of the cD galaxy and the $z=0.9$ lensed system. Throughout the paper we assume $H_{o} = 50$ km s$^{-1}$Mpc$^{-1}$ and $q_{o}$ = 0.5. ",
        "conclusions": "In this letter, we restrict the discussion to the 4 sources found in the maximum sensitivity area of ISOCAM rasters and seen both at 6.7 $\\mu$m and 15 $\\mu$m, i.e. \\#1--4. Their photometric properties are summarized in Table \\ref{sources}. The visible and near-IR spectral energy distributions (SED) computed for these 4 sources are given in Table \\ref{SED}, and compared to those obtained for a typical cluster galaxy. Details on these photometric data can be found in Pell\\'o et al. (1998) and the references therein. In addition, about 20 objects are identified at 6.7 $\\mu$m and 10 at 15 $\\mu$m, which will be discussed in a forthcoming paper. All sources are point-like for ISOCAM, except the cD galaxy which extends over two times the PSF FWHM at 6.7 $\\mu$m (i.e $\\sim 20$ kpc), and then allow us to exclude a pure AGN emission.  \\subsection{The cD galaxy and its cooling flow} The cD galaxy is detected both at 6.7 $\\mu$m and 15 $\\mu$m, with a flux of $300^{+50}_{-40} \\mu$Jy and $500^{+80}_{-70} \\mu$Jy respectively. VLA observations (\\cite{Arnouts}) show a point-like source with decreasing radio fluxes of 140, 16 and 5.5 mJy at 6, 2 and 1.3 cm, respectively. Assuming a power law spectrum, the mid-IR flux would be some $10^5$ fainter than observed, which excludes a jet-like synchrotron contribution to the observed mid-IR emission. In galaxies where the mid-IR emission is dominated by an old stellar population, the ratio 6.7 $\\mu$m/15 $\\mu$m is $> 1$. An excess of 15 $\\mu$m emission in field galaxies indicates the presence of dust heated by UV photons from star-forming regions. Compared to the 6.7 $\\mu$m/15 $\\mu$m ratios observed in other nearby early-type galaxies (\\cite{Suzanne}) or in distant cD galaxies in clusters (\\cite{next}), the ratio of $\\sim 0.6$ found here for the cD is exceptional. This ratio is compatible with the colors of the disk component of Centaurus A (\\cite{Felix}), a nearby giant early-type galaxy exhibiting active star-forming regions in dust lanes, due to a merge with a spiral galaxy. Thus, the cD in A2390 is probably also undergoing active star formation. However, our cD galaxy looks notably different in other wavelengths. Cen A shows a jet plus an extended radio emission but our cD does not. In addition, a B image reveals the existence of a filament extended along the main axis of the cD, while V and I HST images show the presence, within the filament, of very blue globules possibly associated with the 15 $\\mu$m maximum (Fig. \\ref{zoom}). Strong emission lines are present across the long-slit spectrum of the filament (Fig. \\ref{Halpha}) with ratios indicative of massive star formation associated with shocks and incompatible with an active nucleus (\\cite{Allen}, \\cite{Baldwin}). Finally the SED of the filament exhibits a clear excess in the V and B bands with respect to what is expected for a typical elliptical. Assuming that the V flux is mainly produced by forming stars, we derive $M_V = -20.8 \\pm 0.2$ and a SFR of $8 \\pm 4$ h$_{50}^{-2}$ M$_{\\odot}$ yr$^{-1}$ for the optical filament, in agreement with the values obtained from the B flux, corrected for the $[OII]3727$ emission. No absorption has been considered in this calculation, so this estimated SFR has to be taken as a lower limit. According to the results derived from V and B band, and from IR, the optical light is probably coming from the most external regions of the star-forming system, whereas part of the star-formation activity remains shrouded within more dense and dusty clouds, as in the Antennae Galaxies, where absorption is ten times higher when derived from mid-IR than from J, H, K bands, as most massives stars are not visible at optical wavelengths (\\cite{antenne}). This implies that we can not exclude a SFR as high as ten times what is derived from the optical for the cD of A2390. Those differences explain why, with a 6.7 $\\mu$m/15 $\\mu$m ratio of 0.6, we can not just consider the cD of A2390 as an early-type galaxy undergoing simple star-formation in dust lanes as Cen A, but most probably as the place of one or several massive star-bursts which may be located in the central globules (see Fig \\ref{zoom}). Indeed, the study of the X-ray image of this cluster demonstrated the presence of one of the strongest cooling flows known ($\\sim 800$ M$_{\\odot}$ yr$^{-1}$ within a cooling radius of 200 kpc), surrounding the cD galaxy (\\cite{marguerite}). Giving the size of the mid-IR emitting region in the cD, $\\sim 20$ kpc, we derive a mass flow of 80 M$_{\\odot}$ yr$^{-1}$, assuming that $\\dot{M} \\sim r$ (\\cite{fabian}). Note that 20 kpc is also about the size of the optical filament (Fig. \\ref{zoom}). However, our present understanding of the relationship between mid-IR dust emission and star formation is still too preliminary to infer quantitative constraints on the IMF or even on the heating processes involved in this complex medium. Finally, the total star-formation rate of $\\sim 10$ M$_{\\odot}$yr$^{-1}$ deduced from the optical in the filament is clearly a lower limit, and the huge quantity of gas needed could be provided by the cooling flow. However, despite the fact that spiral galaxies are very rare in the core of rich clusters, the hypothesis of a past merge with a late-type galaxy cannot be formally excluded here, which would also provide gas for some $10^{7-8}$ years. \\begin{figure} \\psfig{file=lemonon.center3.ps,width=8.5cm} \\caption []{The V-I color HST image, obtained from F555W and F814W, of the core of the cD galaxy of A2390. White and black crosses indicate 6.7 $\\mu$m and 15 $\\mu$m maximum emission respectively with the associated uncertainties (see text). Some very blue globules suspected to be active star forming regions could be responsible for the 15 $\\mu$m emission which coincides with the most luminous one, and associated with strong cooling flow ($\\sim800$ M$_{\\odot}$yr$^{-1}$) surrounding this cD galaxy. They are part of a NW--SE filament, aligned with the cD main axis, and the overall cluster X-ray elongation on large scale.} \\label{zoom} \\end{figure} \\begin{figure} \\psfig{file=lemonon.halpha.ps,width=9.8cm} \\caption []{Mean long-slit spectrum of the filament within the cD galaxy (without reddening correction), from CFHT (\\cite{Arnouts}).} \\label{Halpha} \\end{figure} \\subsection{What is new in the arc system of Abell 2390 ?} After the detection of the giant arc at $z=0.724$ in Abell 370 (\\cite{Metcalfe}), observation of the complex arc system of Abell 2390 confirms the capability of ISOCAM to point up very distant lensed objects. The giant arc consists of three parts, A at $z=1.033$ (\\cite{brenda}), and B--C, at $z=0.913$ (\\cite{Pello}). Near IR imaging already distinguished A from B--C, as A was not detected in the K band (\\cite{Smail}). HST images revealed that B and C are likely two interacting galaxies. The present ISOCAM images are in full agreement with this picture. Although it was not possible to estimate properly the 6.7 $\\mu$m flux because of blending, the 15 $\\mu$m/6.7 $\\mu$m ratio for the B--C component is well larger than unity which is indicative of the presence of an active star forming region in agreement with the strong [OII] line detected in the optical spectrum (\\cite{Pello}). Except for its lower amplification factor, the case of object D is very similar. Its morphology in the HST images is complex with probable signs of interaction and low surface brightness extensions. The existence of starburts in the two interacting galaxies is then not a surprise. The optical and near-IR SEDs of objects B--C and D appear brighter in the near-IR and fainter in the blue bands compared to A. These SEDs can be fitted by different synthetic spectra at $z = 0.913$, using the GISSEL96 code (Bruzual \\& Charlot, 1998) to approximately constrain the parameters, and a single stellar population (instantaneous burst), an extinction curve of SMC type (Pr\\'evot et al. 1984), and assuming the Scalo IMF (1986). The best fits of the sources B--C and D are obtained with a rest-frame $A_{V} \\sim3$ in both cases, a stable result with respect to metallicity changes. The corrected magnitudes for objects B--C and D (lensing and absorption) are very similar ($M_B = -20.8$) , the total mass involved in the burst being $\\sim10^{10} M_{\\odot}$ in both cases. Despite uncertainties on burst age, a constant star-forming model gives similar results and a mean corrected SFR of $40$ to $50\\, M_{\\odot} h_{50}^{-2} yr^{-1}$. According to these results, the two lensed sources detected by ISOCAM at $z=0.913$ are strongly reddened star-forming galaxies. In the case of A, there is no need for a reddening correction to fit the SED.  Finally, the ISOCAM source \\#4 detected in both channels may be associated with a very faint source in the HST image (I = 23.5), with a fuzzy shape. Its 15 $\\mu$m/6.7 $\\mu$m ratio is very high ($\\sim 3$). A photometric redshift of $z=0.4^{+0.2}_{-0.08}$ is proposed for this object by techniques described by Miralles \\& Pell\\`o (1998). Even if the results are much more uncertain in this case ($75\\%$ confidence), the best fit of the SED gives $A_{V} \\sim$3.5--4.2 in order to explain the high J and K$^\\prime$ emission compared to the optical bands. The corrected SFR is relatively low, $\\sim1 M_{\\odot} h_{50}^{-2} yr^{-1}$. Taking the photometric redshift into account, the SED of this object, with strong mid-IR emission with respect to its optical counterpart, is probably dominated by the so-called unidentified infrared band emitters, and its colors are similar to those of the post-starburst companion of M51 (Boulade et al., 1996).  \\subsection{Summary and conclusion}  From deep and high-resolution ISOCAM images of the core of Abell 2390 we discovered active star forming regions in the two most distant lensed galaxies ever seen in a cluster by ISO. This allowed us to support the scenario of two interacting galaxies at $z=0.913$ in the ``straight arc'' of A2390, as well as in the other galaxy observed at the same redshift. More interesting, we detect a very faint emission from the cD galaxy at 6.75 $\\mu$m, compared to other cluster dominant galaxies at similar redshift (\\cite{next}). But, for the first time, the strong 15 $\\mu$m/6.75 $\\mu$m emission ratio flags the presence of a large amount of warm dust in the cD, probably associated with a very active star forming region located within the envelope of the galaxy. Thus, our observation may further elucidate the fate of part of accumulating gas in the complex cooling flow radio core environment."
    },
    "9804/astro-ph9804277_arXiv.txt": {
        "abstract": "Using an Eulerian perturbative calculation, we show that the distribution of relative pairwise velocities which arises from gravitational instability of Gaussian density fluctuations has asymmetric (skewed) exponential tails. The negative skewness is induced by the negative mean streaming velocity of pairs (the infall prevails over expansion), while the exponential tails arise because the relative pairwise velocity is a {\\it number}, not volume weighted statistic. The derived probability distribution is compared with N-body simulations and shown to provide a reasonable fit. ",
        "introduction": "\\label{sec-intro}  Redshift surveys present a distorted picture of the world because peculiar motions displace galaxies from their true spatial positions. This phenomenon, which would make redshift surveys useless for intergalactic spaceship navigators, is extremely useful for cosmologists. It can serve as a probe of the dynamics of gravitational clustering and the cosmological mass density parameter, $\\Omega$ (\\cite{sar77}; \\cite{pee80}, hereafter LSS; \\cite{kai87}; \\cite{ham92}; \\cite{pee93}, hereafter PPC; \\cite{reg95}). A convenient statistical measure of the distortion effect is the galaxy two-point correlation function in redshift space. Under certain assumptions it can be expressed as a convolution of the true spatial correlation function, $\\xi(r)$, with the distribution of the relative line-of-sight velocities of pairs of galaxies, $p(w|r,\\theta)$ . Here $r$ and $w$ are respectively, the spatial separation and relative radial velocity of a pair of galaxies, while $\\theta$ is the angle between the separation vector ${\\bf r}$ and observer's line of sight (cf. LSS; \\cite{fis95}, hereafter F95). The purpose of this {\\it Letter} is to derive $p(w|r,\\theta)$, using weakly nonlinear gravitational instability theory. This distribution was measured from N-body simulations and estimated indirectly from redshift surveys. At $r \\simlt$  1\\mpc, where\\footnote{We use the standard parametrisation for the Hubble constant, $H = 100 \\, h$ \\kms Mpc$^{-1}$.} the galaxies are strongly clustered ($\\xi \\simgt 20$), the observations are consistent with an exponential distribution (\\cite{pee76}; \\cite{dav83}; \\cite{fis94}, hereafter F94; \\cite{mar95}; \\cite{lsd97}). The fact that $p(w)$ at small separations differs strongly from its initial, Gaussian character, is not surprising: after all, the small-scale velocity field has been `processed' by strongly-nonlinear dynamics in clusters, and exponential distributions were recently derived from the \\cite{pre74} theory (\\cite{she96}, \\cite{dia96}). On larger scales, where the fluctuations have small amplitudes, one na{\\\"\\i}vely expects to see the `unprocessed' initial conditions. However, N-body experiments suggest that $p(w|r,\\theta)$ retains its exponential character even at separations $r \\simgt$ 10\\mpc, where $\\xi \\simlt 0.1$, despite the fact that the initial density and velocity fields in those experiments were drawn from a Gaussian distribution (\\cite{efs88}, hereafter EFWD; \\cite{zur94}, hereafter ZQSW; F94). At similar separations, an exponential $p(w|r,\\theta)$ has also been inferred from observations (F94; \\cite{lov96}). The simulations also show that the radial component of the distribution, $p(w|r,0^{\\circ})$ is significantly skewed, in particular at large separations (EFWD; ZQSW; F94). The physical origin of the skewness and exponential shape of $p(w)$ at large separations has until now remained unexplained. We provide the explanation below. ",
        "conclusions": ""
    },
    "9804/astro-ph9804310_arXiv.txt": {
        "abstract": "We have examined the hypothesis that the majority of the diffuse EUV flux in the Coma cluster is due to inverse Compton scattering of low energy cosmic ray electrons ($0.16 < \\epsilon < 0.31$\\,GeV) against the 3$^{\\circ}$\\,K black-body background. We present data on the two-dimensional spatial distribution of the EUV flux and show that these data provide strong support for a non-thermal origin for the EUV flux. However, we show that this emission cannot be produced by an extrapolation to lower energies of the observed synchrotron radio emitting electrons and an additional component of low energy cosmic ray  electrons is required. ",
        "introduction": "Diffuse EUV emission has been detected in five clusters of galaxies: Virgo (Lieu et al., 1996a), Coma (Lieu et al., 1996b), Abell 1795 (Mittaz et al., 1997), and Abell 2199 and Abell 4038 (Bowyer et al., 1997). These clusters were detected with a statistical significance varying from 8 to 50 standard deviations. The diameter of the diffuse emission in these clusters ranges from 20$\\arcmin$ to 40$\\arcmin$.  Some diffuse EUV emission in clusters of galaxies would be expected from the well-studied X-ray cluster emission. However, in all cases examined to date, the EUV emission is far greater than the expected emission from the X-ray-emitting gas. Marginal signatures of this ``soft excess'' are sometimes present in the lowest energy resolution band of the ROSAT PSPC, where they produce less than a 20\\% enhancement over the emission expected from the X-ray cluster gas. In contrast, the excesses found with EUVE range from 70\\% to 600\\% above that expected from the X-ray gas. A variety of instrumental and Galactic interstellar medium absorption effects have been suggested as alternative explanations for the EUVE data, but these have all been found wanting (For a discussion of these issues, see Bowyer et al., 1997).  In the original reports of diffuse EUV cluster emission, the data were interpreted in terms of additional thermal gas components in the clusters. In these analyses, the known X-ray emission from the hot cluster gas was first fitted to the EUV data, and the excess EUV emission was computed. This excess was then fitted by additional components of thermal gas. Because of the low energy of the EUV emission, much lower temperature thermal gases ($\\sim 10^6$\\,K) are required. The concept that additional components of lower temperature gas are present in these clusters has not received wide support, primarily because gas at these temperatures is near the peak of the radiative cooling curve and hence cools rapidly, requiring a substantial energy input to sustain the gas at these temperatures. In addition, it is difficult to understand how different components of gas at grossly different temperatures could retain their separate identity. For example, the Coma cluster has been shown to be formed by merging of distinct subunits, in both X-ray (White et al., 1993), and optical (Colless and Dunn, 1996) studies. This produces variations less than a factor of 3 in temperatures of the X-ray emitting gas (Honda et al. 1996). Deiss and Just (1996) have shown in a general analysis, and with a specific application to the Coma cluster, that turbulent mixing time scales are only a few $10^9$\\,years, which argues against the co-existence of major quantities of gas at two vastly different temperatures. However, Cen et al. (1995) have argued that a warm ($\\simeq 10^6$\\,K) thermal gas is widely distributed throughout the universe, as a direct product of the growth of structure leading, eventually, to clusters of galaxies. In their scenario, the energy required to sustain the warm gas is provided by gravitation.  Hwang (1997) has examined the hypothesis that the source of the diffuse EUV flux is inverse Compton scattering by electrons that are a low energy extrapolation of electrons producing the observed synchrotron emission; these electrons are scattered  against the 3$^{\\circ}$\\,K black-body background radiation. The magnetic field he derived for the cluster was 0.2 to 0.4\\,$\\mu$G which is consistent with the range of estimates for the cluster field. In this work he only considered the constraints imposed by the total EUV flux.  En{\\ss}lin and Biermann (1998) also considered this mechanism as the source of the EUV flux in the Coma cluster. They assumed that the relativistic energy density of the synchrotron emitting electrons scale radially with the same profile as the X-ray producing gas. This assumption can be questioned since the non-thermal relativistic electrons may well be independent of the thermal X-ray gas. These authors cite the best support for this assumption is given by the data in Figure\\ 3 of Deiss et al. (1997) which compares the X-ray and radio radial emission profiles. Unfortunately, the results in this figure are incorrect as has been confirmed by Deiss (private communication). These authors find a magnetic field of 1.2\\,$\\mu$G is required given their assumptions; this field is also consistent with the range of estimates for the cluster.  Sarazin and Lieu (1998) have explored the possibility that EUV radiation in clusters of galaxies could be produced by inverse Compton scattering by a population of very low energy cosmic ray electrons. They showed that the one dimensional EUV spatial profile for the cluster A1795, a radio quiet cluster, was consistent with this hypothesis. A potential problem with this hypothesis as a universal explanation for the EUV emission in clusters of galaxies is that the electrons they proposed have an energy density and pressure which are 1 to 10\\% of that of the thermal gas in clusters. If one includes the pressure of cosmic ray ions to the pressure of the electrons proposed by Sarazin and Lieu, using the ratio expected on theoretical grounds (Bell 1978) and measured at Earth orbit (see. e.g., Weber 1983), the total cosmic ray pressure is substantially larger than that of the X-ray emitting gas.  In this work we reconsider the hypothesis that the EUV emission in the Coma cluster is the result of inverse Compton emission. We first consider the constrains imposed by the total EUV flux. We review the existing radio data and obtain a different spectral index than that employed by Hwang (1997) and En{\\ss}lin and Biermann (1998), which we argue is more appropriate. We then derive results which are generally consistent with the inverse Compton hypothesis. We then consider the two dimensional spatial distribution of the EUV flux; this data provides substantial support for a non-thermal origin for this flux. We find that the spatial distribution of the magnetic field required by En{\\ss}lin and Biermann to produce the EUV emission profile is unrealistic. We show that the difference in spatial extent between the EUV and radio halos cannot be explained using an electron distribution which is an extrapolation of the known synchrotron emitting electrons and that an additional population of low energy cosmic ray electrons are required to explain these data. ",
        "conclusions": "We have examined the hypothesis that the EUV radiation from the Coma cluster is due to inverse Compton scattering of low energy cosmic ray electrons against the 3$^{\\circ}$\\,K black-body background radiation. The total integrated EUV emission produced by cosmic ray electrons which are a low energy extrapolation of higher energy electrons, known to be present from their synchrotron emission, gives results which are consistent with the range of estimates of the magnetic field in the cluster.  We next consider the two dimensional spatial distribution of the EUV emission. This emission does not follow the distribution of the gravitationally bound X-ray gas, but rather exhibits an asymmetric distribution similar to that exhibited by the radio emission. This suggests a non-thermal origin for the EUV emission rather than a gravitationally constrained thermal gas.  We show from a comparison of the size of the EUV halo and the radio halo that the EUV emission cannot be produced by inverse Compton radiation from electrons which are an extrapolation of the distribution which produces the observed radio emission. We develop a model for the EUV emission which is self-consistent and fits the existing data. This model requires an additional component of low energy cosmic rays.  Inverse Compton EUV emission is surely present at some level in clusters of galaxies with radio halos. However, it may well be masked by emission from some other more dominant source mechanism. A test of the inverse Compton scattering hypothesis as the source of the EUV flux in the Coma cluster would be provided by a measurement of the size of the radio halo at $\\sim 1$\\,MHz. Unfortunately, this is a challenging measurement because of instrumental limitations and ionospheric effects. In addition, at these low frequencies self-absorption could affect the surface brightness profile by reducing the flux near the cluster center while increasing the halo size; this effect would have to carefully be considered when interpreting such a measurement."
    },
    "9804/astro-ph9804126_arXiv.txt": {
        "abstract": "Cosmological implications of clusters of galaxies are discussed with particular attention to their importance in probing the cosmological parameters. More specifically we compute the number counts of clusters of galaxies, \\ns relation, in X-ray and submm bands on the basis of the Press--Schechter theory. We pay particular attention to a set of theoretical models which well reproduce the {\\it ROSAT} 0.5-2 keV band \\ns, and explore possibilities to break the degeneracy among the viable cosmological models. ",
        "introduction": "There are several reasons why clusters of galaxies are regarded as useful probes of cosmology including (i) since dynamical time-scale of clusters is comparable to the age of the universe, they should retain the cosmological initial condition fairly faithfully.  (ii) clusters can be observed in various bands including optical, X-ray, radio, mm and submm bands, and in fact recent and future big projects (e.g., SDSS, AXAF, PLANCK) aim to make extensive surveys and detailed imaging/spectroscopic observations of clusters.  (iii) to the first order, clusters are well approximated as a system of dark matter, gas and galaxies, and thus theoretically well-defined and relatively well-understood, at least compared with galaxies themselves, and (iv) on average one can observe a higher-z universe with clusters than with galaxies.  In particular X-ray observations are well-suited for the study of clusters since the X-ray emissivity is proportional to $n_e^2$ and thus less sensitive to the projection contamination which has been known to be a serious problem in their identifications with the optical data.  In fact, various statistics related to the abundances of clusters has been extensively studied to constrain theories of structure formation, including mass function (Bahcall \\& Cen 1993; Ueda, Itoh, \\& Suto 1993), velocity function (Shimasaku 1993; Ueda, Shimasaku, Suginohara, \\& Suto 1994), X-ray Temperature function (hereafter XTF, Henry \\& Arnaud 1991; White, Efstathiou, \\& Frenk 1993; Kitayama \\& Suto 1996 ; Viana \\& Liddle 1996; Eke, Cole, \\& Frenk 1996; Pen 1996).  Previous authors have focused on cosmological implications of cluster XTF mainly because theoretical predictions are relatively easier although the observational data are statistically limited. In addition, the conversion to the number density at high z becomes very sensitive to the adopted cosmological parameters.  On the other hand, \\ns which we discuss in details below is observationally more robust (Ebeling et al. 1997; Rosati \\& Della Ceca 1997) while its theoretical prediction is more model-dependent (Oukbir, Bartlett, \\& Blanchard 1996; Kitayama \\& Suto 1997; Kitayama, Sasaki \\& Suto 1998).  In this respect, both statistics are complementary.  \\begin{figure}[t] \\begin{center} \\leavevmode \\psfig{figure=logns1.cps,height=6.8cm} \\psfig{figure=chi2ns2.cps,height=6.8cm} \\end{center} \\vspace*{-0.5cm} \\caption{ {\\it Left:} Theoretical predictions for \\ns of X-ray clusters in CDM models with different cosmological parameters; (a) $\\sigma_8=1.04$ models with different $\\Omega_0$, $\\lambda_0$ and $h$, (b) $\\Omega_0=1$ and $0.45$ models with different $\\sigma_8$. Denoted by (COBE) are the models normalized according to the {\\it COBE} 4 year data (Bunn \\& White 1997).  Data points with error bars at $S\\simlt 10^{-12}$ \\unit are from the {\\it ROSAT} Deep Cluster Survey (RDCS, Rosati et al. 1995; Rosati \\& Della Ceca 1997), and the error box at $S\\simgt 2 \\times 10^{-12}$ represents a power-law fitted region from the {\\it ROSAT} Brightest Cluster Sample (BCS, Ebeling et al. 1997). For the BCS data at $S= 2 \\times 10^{-12}$, $1 \\times 10^{-11}$ and $6 \\times 10^{-11}$\\unit, we also plot the corresponding Poisson errors.  {\\it Right:} Limits on $\\Omega_0$ and $\\sigma_8$ in CDM models ($n=1$, $h=0.7$) with (a) $\\lambda_0=1-\\Omega_0$, and (b) $\\lambda_0=0$. Constraints from cluster \\ns (solid) and XTF (dotted) are plotted as contours at $1 \\sigma$(68\\%), $2\\sigma$(95\\%) and $3\\sigma$(99.7\\%) confidence levels. Dashed lines indicate the {\\it COBE} 4 year results from Bunn \\& White (1997).  } \\label{fig:ns1chi2} \\vspace*{-0.5cm} \\end{figure} ",
        "conclusions": "Let us summarize the conclusions of the present talk.  \\begin{description} \\item{(1)} There exist { several theoretical models} which successfully reproduce the observed \\ns relation of galaxy clusters over almost four orders of magnitude in X-ray flux.  \\item{(2)} The resulting $\\sigma_8$ is given by the following empirical fit (95\\% confidence limit): \\begin{equation} \\sigma_8 = (0.54 \\pm 0.02 \\pm 0.1) \\times \\Omega_0^{-0.35-0.80\\Omega_0+0.55\\Omega_0^2} \\end{equation} for $\\lambda_0=1-\\Omega_0$ CDM, and \\begin{equation} \\sigma_8 = (0.54 \\pm 0.02 \\pm 0.1) \\times \\Omega_0^{-0.28-0.91\\Omega_0+0.68\\Omega_0^2} \\end{equation} for $\\lambda_0=0$ CDM.  \\item{(3)} Low-density CDM models ($n=1$) with $(\\Omega_0,\\lambda_0,h,\\sigma_8) = (0.3,0.7,0.7,1)$ and $(0.45, 0, 0.7, 0.8)$ simultaneously account for the cluster \\ns, XTF, the {\\it COBE} 4 year normalization. \\end{description}   Maybe the most important point is that many cosmological models are more or less successful in reproducing the structure at redshift $z\\sim0$ {\\it by construction}. This is because the models have still several degrees of freedom or {\\it cosmological parameters} which can be appropriately {\\it adjusted} to the observations at $z\\sim0$ ($\\Omega_0$, $\\sigma_8$, $h$, $\\lambda_0$, $b(r,z)$).  We have shown that, given a complete flux limited cluster sample with redshift and/or temperature information, one can further constrain the cosmological models. In fact, our tentative comparison indicates that our predictions reproduce well the evolutionary features of the XBACs and that the results, although preliminary, seem to favor low density ($\\Omega_0 \\sim 0.3$) universes.  As indicated by this preliminary result, surveys of objects at at high redshifts in several different bands (X-ray, radio and submm) are the most efficient and rewarding to break the degeneracy among the viable cosmological models."
    },
    "9804/astro-ph9804256_arXiv.txt": {
        "abstract": "Steep soft X-ray (0.1-2 keV) quasars share several unusual properties: narrow Balmer lines, strong FeII emission, large and fast X-ray variability, rather steep 2-10 keV spectrum. These intriguing objects have been suggested to be the analogs of Galactic black hole candidates in the high, soft state. We present here results from ASCA observations for two of these quasars: NAB0205+024 and PG1244+026.  Both objects show similar variations (factor of $\\sim2$ in 10 ks), despite a factor of about ten difference in the 0.5-10 keV luminosity ($7.3\\times 10^{43}$ erg s$^{-1}$ for PG1244+026 and $6.4\\times10^{44}$ erg s$^{-1}$ for NAB0205+024, assuming isotropic emission, $H_0$ = 50.0 and $q_0$ = 0.0).  The X-ray continuum of the two quasars flattens by 0.5-1 going from the 0.1-2 keV band toward higher energies, strengthening recent results on another half dozen steep soft X-ray AGN.  PG1244+026 shows a significant feature in the `1 keV' region, which can be described by either as a broad emission line centered at 0.95 keV (quasar frame) or as edge or line absorption at 1.17 (1.22) keV.  The line emission could be due to reflection from an highly ionized accretion disk, in line with the view that steep soft X-ray quasars are emitting close to the Eddington luminosity.  Photoelectric edge absorption or resonant line absorption could be produced by gas outflowing at a large velocity (0.3-0.6 c). ",
        "introduction": "The ROSAT PSPC has found a large spread in the energy spectral indices of low-z quasars\\footnote{We use ``quasars'' to describe broad line emission objects, regardless of luminosity.} : $0.5<\\alpha_{0.1-2keV}<3.5$. In about 10$\\%$ of cases $\\alpha_{0.1-2keV}\\gs 2$ (e.g. Laor et al. 1994, 1997, Walter \\& Fink 1993, Fiore et al. 1994). The large spread in $\\alpha_{0.1-2keV}$ favoured the discovery of its correlation with other properties. In fact, the steep soft X-ray quasars have then been realized to share a cluster of unusual properties:  \\begin{itemize}  \\item narrow Balmer lines \\footnote{the permitted lines have FWHM$\\ls$2000~km~s$^{-1}$, yet still are clearly broader than the forbidden lines.} ( Laor et al 1994, 1997, Boller et al. 1995);  \\item strong FeII emission (Laor et al 1994, 1997, Lawrence et al. 1997)  \\item Rapid, large amplitude variability (factor of 2-50 on timescales from minutes to months, Boller et al., 1995, Brandt et al., 1995, Otani 1995, Boller et al. 1997)  \\item Somewhat steep hard X-ray spectra (2$>\\alpha_{2-10keV} >0.6$, Pounds et al. 1995, Brandt et al., 1997);  \\end{itemize}  Pounds et al. (1995), suggest the latter to be a close physical analogy with the X-ray power-law produced by Comptonization in a hot accretion disk corona in Galactic black hole candidates (BHC) in their `soft-high' state. This is not the only analogy between BHC and steep X-ray spectrum quasars.  Laor et al. (1994, 1997) explained the correlation with H$\\beta$ FWHM as due to the larger size of a virialized broad emission line region for an AGN in a high $L/L_{Edd}$ state.  Ebisawa (1991) found that while the soft component of 6 BHC observed by Ginga is roughly stable on time scales of 1 day or less, the hard component exhibits large variations down to msec time scales. These timescales translates to $10^4$ years and 0.1 day for quasars, if they scale with the mass of the compact object.  The soft component of BHC extends up to $\\sim 10$ keV in BHC in `soft-high' states, and it is often associated with optically thick emission from an accretion disk. If this is the case, the temperature should scale with the mass of the compact object as $M_{BH}^{-1/4}$, and the above energy translates to 0.1-0.4 keV for quasars.  The rapid large amplitude variability shown by a few narrow line Seyfert 1 galaxies (NLSy1) at about 1 keV on timescales of hours to days (Otani 1995, Brandt et al. 1995, Boller et al. 1997) can then be analogous to the above BHC hard component flickering.  A steep X-ray spectrum quasar with 10-100 times the luminosity of NLSy1s, should be larger and so should vary no more rapidly than several days. Instead Fiore et al. (1998a) find that steep X-ray spectrum PG quasars commonly vary by a factor 2 in 1~day. Variability seems therefore correlated with X-ray spectral slope and Balmer line width (and therefore possibly with the accretion rate) rather than with the luminosity.  Evidence for spectral features in the `1 keV' region in many steep soft X-ray quasars is building up (Turner et al., 1991, Brandt et al., 1994, Otani et al., 1995, Comastri et al., 1995, Leighly et al., 1997, 1998a,b).  Instead, `normal' Seyfert 1 galaxies (having broad Balmer lines and flatter soft X-ray spectra) usually have their strongest absorption features at lower energies (in the 0.6-0.9 keV `oxygen' band).  An intriguing possibility is that the appearance of these features at different energies also depend on $L/L_{\\rm Edd}$.  Detailed high energy X-ray spectra of luminous quasars with steep soft X-ray spectra are essential to understand the `narrow-broad line' phenomenon in AGN, in particular whether the peculiar X-ray properties depend on optical luminosity, optical-to-X-ray ratio ($\\alpha_{OX}$), or on their Eddington ratio.  To this end we selected two bright quasars with $\\alpha_{0.1-2keV}>$2.0 (Fiore et al., 1994) at the extreme values of optical luminosity, both with low Galactic $N_H$ (Table 1) of $1.9\\times 10^{20}$ cm$^{-2}$ for PG~1244+026, and of $3.0\\times 10^{20}$ cm$^{-2}$ for NAB0205+024, Elvis et al., 1989) and observed them with ASCA. We report the results in this paper. ",
        "conclusions": "ASCA observations of two steep soft X-ray quasars have shown that:  \\begin {enumerate}  \\item The X-ray continuum of the two quasars flattens by $\\Delta\\alpha=0.5-1$ going toward high energies. Similar results were obtained by authors on some half dozen steep soft X-ray quasars.  \\item PG1244+026 shows a significant feature in the `1 keV' region.  Similar features were again reported in other steep soft X-ray quasars. The data are not good enough to discriminate between a broad emission line centered at 0.95 keV (quasar frame) or an absorption edge at 1.17 keV, or an absorption notch at 1.22 keV.  Line emission could be due to reflection from an highly ionized accretion disk, in line with the view that steep soft X-ray quasars are emitting close to the Eddington luminosity.  Photoelectric edge absorption or resonant line absorption could be produced by gas outflowing at a large velocity (0.3-0.6 c). In these absorption models significant cold (i.e. oxygen less ionized than OVI) absorption in excess of the Galactic is required. This would imply an increase by a factor 2-3 of the cold column with respect to a previous PSPC observation or a peculiar ionization structure.  In neither the emission or absorption cases the SIS resolution is good enough to identify unambiguously the ions responsible for the feature. The high resolution and high throughput of the low energy gratings and spectrometers of AXAF and XMM are clearly needed to shed light on this puzzling case.  \\item The two quasars show similar variability properties (flux variations up to a factor of 2 in 10 ks) despite a factor of ten difference in the X-ray observed luminosity. This agrees with the Fiore et al (1998a) finding that the variability properties of radio-quiet quasars are correlated with the shape of the X-ray spectrum, the width of the Balmer lines and so possibly with the accretion rate.   \\end{enumerate}  \\bigskip F.F acknowledges support from NASA grants NAG 5-2476 and NAG 5-3039, B.J.W. acknowledges support from ASC contract NAS8-39073."
    },
    "9804/astro-ph9804060_arXiv.txt": {
        "abstract": "Since it has become possible to discovery planets orbiting nearby solar-type stars through very precise Doppler-shift measurements, the role of methods used to analyze such observations has grown significantly. The widely employed model-dependent approach based on the least-squares fit of the Keplerian motion to the radial-velocity variations can be, as we show, unsatisfactory. Thus, in this paper, we propose a new method that may be easily and successfully applied to the Doppler-shift measurements. This method allows us to analyze the data without assuming any specific model and yet to extract all significant features of the observations. This very simple idea, based on the subsequent subtraction of all harmonic components from the data, can be easily  implemented. We show that our method can be used to analyze real 16 Cygni B Doppler-shift observations with a surprising but correct result which is substantially different from that based on the least-squares fit of a Keplerian orbit. Namely, using  frequency analysis we show that with the current accuracy of this star's observations it is not possible to determine the value of the orbital eccentricity which is claimed to be as high as 0.6. ",
        "introduction": "Recent improvements in the long-term precision of Doppler-shift measurements \\cite[]{Marcy:96c::} resulted in several spectacular detections of planetary companions to solar-type stars \\cite[for review see the paper of][] {Marcy:97::}. As such discoveries supply indirect evidence of the existence of extra-solar planets,  other explanations of observed radial-velocity variations appeared, e.g., stellar pulsations \\cite[]{Gray:97::}. The most recent results, however, show that only a planetary hypothesis is acceptable \\cite[]{Marcy:98::,Gray:98::}.  The usual procedure showing that there exists a planet around s star consits of a direct least-square fit of Keplerian model to  the observations. It always gives   certain values of orbital parameters and their formal errors. In the case of  `good' data this is the best and the quickest way to obtain relable results. However, in the case of spare data with big errors one has to prove  that  the least-squares method can be used  and that the obtained  parameter values   and their errors are good estimates of the real values. This is a difficult  and time consuming task. Without doubt, we have this situation  with the  Doppler observation of extra-solar planets.  We present an analysis of this problem and  we show that the eccentricity of the fitted orbit is a very sensitive parameter and, in some cases, its value   and   error  given by the least-squares method are not correct.  The aim of this paper is to show how the mentioned problem can be solved in practice. Namely, we propose a method  that can be very useful for analyzing radial-velocity variations.  It is based on a simple idea involving the subsequent subtraction of periodic components from the data. This approach allows us to analyze the observations without assuming any specific model describing the system  behavior (like Keplerian motion or stellar pulsations). After the determination of all significant components of the data, it remains to be decided which process is responsible for what we observe and whether it is possible to choose only one.  The plan of this paper is  as follows. In Section 2 we  analyse observations  of 16 Cygni B and we explain why  the standard least-squere fit does not give reliable estimates of parameters and their errors. In Section 3 we analyze analytically the  Keplerian motion of the system `a star with one planet' in order to learn how its motion modulates the observed star radial-velocities. We investigate  mainly the spectral properties of the motion which are essential for our method. In Section 4 we develop a simple numerical technique which can be used to extract all the information we need to compare with the results from Section 3. In Section 5, we perform a numerical test of the method using  simulated radial-velocity variations with the orbital parameters of 70 Vir  \\cite[]{Marcy:96b::}. In section 6, we discuss the application of the method to finding the eccentricity of 16 Cygni B \\cite[]{Cochran:97::}. ",
        "conclusions": "In this paper we have shown that results obtained from least-squares fits of Keplerian orbits to real Doppler-shift measurements may lead to incorrect interpretations. Specifically, they may give unrealistic or even entirely false values of parameters and their uncertainities. In order to solve these problems we have proposed a new method, frequency analysis, which efficiently provides an independent test of the reliability of determined orbital parameters. This method may deliver a substantial revision of the current values of planets' high eccentricities that are essential for our understanding of the formation and evolution of planetary systems. It might even lead to hints that some of the observed high eccentric planets are in fact planetary systems consisting of more than one planet or at least provide an independent point of view on the same data. These facts, together with the ease of applicability of frequency analysis, make our method worth trying on future observations if not for the data already gathered."
    },
    "9804/physics9804027_arXiv.txt": {
        "abstract": "It is shown that the plasma, generated during an impact of a meteoroid with an artificial satellite, can produce electromagnetic radiation below the microwave frequency range. This interference is shown to exceed local noise sources and might disturb regular satellite operations. ",
        "introduction": "At the end of 1998, the first modules of the \\emph{International Space Station} will be put in orbit around the Earth and this should open new frontiers for life in space.  The intensive use of the space makes necessary to know the potential risks.  The threat from meteoroids is today well known and several authors have underlined the risks connected with the impact on a spacecraft (for a review, see \\cite{REV}). However, the \\emph{Olympus} end-of-life anomaly \\cite{CASWELL} and the recent work of McDonnell \\emph{et al.} \\cite{MCDONNELL} put a new light on these issues.  The \\emph{Olympus} failure is a paradigmatic example: in that case, the impact with a Perseid meteoroid may have generated electrical failures, leading to a chain reaction which culminated with an early end of the mission \\cite{CASWELL}.  On the other hand, McDonnell \\emph{et al.} \\cite{MCDONNELL} showed that, if the plasma charge and current production during an impact are considered, meteoroid streams can be very dangerous, even during normal conditions.  It should be noted that they considered only damages by direct discharges or current injection in circuits (\\emph{e.g.} via the umbilical) \\cite{MCDONNELL}.  However, there are several other ways by which the plasma could interact with the spacecraft electronics.  For example, it is useful to recall the work of Cerroni and Martelli \\cite{CERRONI}, in which they showed that thermal forces in impact-produced plasmas could explain the magnetisation observed in the neighbourhood of lunar craters.  Even if Cerroni and Martelli studied experimentally hypervelocity impacts of aluminium projectiles on basalt targets, it is possible to extend their work to general hypervelocity impacts.  Here, we show that a plasma cloud, generated during a hypervelocity impact of a meteoroid with an artificial satellite, can radiate electromagnetic energy below the microwave frequency range and, therefore, may disturb regular satellite operations. ",
        "conclusions": "After the \\emph{Olympus} end-of-life anomaly \\cite{CASWELL} and the work of McDonnell \\emph{et al.} \\cite{MCDONNELL}, it seems clear that the meteoroids hazard is not restricted to a mechanical damage. Here it is suggested a new interference path, that is electromagnetic radiation emitted from the impact-produced plasma cloud.  Even if the radiated power is not sufficient to destroy anything, it may disturb regular satellite operations.  Further investigations should be made on specific satellite, because they require detailed information about onboard electronics, in order to calculate possible couplings and non-linearities.  \\stars Author wishes to thank Paolo Farinella, of Department of Mathematics of the University of Pisa, for constructive review. \\newpage"
    },
    "9804/astro-ph9804195_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Hubble Deep Field (HDF)  provides accurate multi-band photometry of galaxies to very faint magnitudes, with 10$\\sigma$ $AB$ magnitude limit of $m_{814}=27.60$ (\\cite{HDF}). The faint limit of the HDF makes it difficult to obtain spectroscopic redshifts for the majority of the galaxies in the field. It is therefore useful to derive estimated redshifts of these galaxies using the available multi-band photometry. Several groups (\\cite{Lanze96}, \\cite{GwHa96}, \\cite{Saw97}) obtained photometric redshifts for the HDF by comparing the observed $UBVI$ fluxes of each object with a set of galaxy spectral templates of different galaxy types redshifted to evenly spaced redshifts. Since spectroscopic redshifts have been measured and published for $\\sim$ 100 galaxies in the HDF (\\cite{Cohen96}, \\cite{Hogg98}; \\cite{Stei}; \\cite{DEEP}), it is possible to fit analytic expressions for photometric redshifts. In this paper, we explore a simple empirical approach to estimating redshifts of galaxies based on their colors (see \\cite{Conno97}, \\cite{Brun97}, \\cite{Conno95} for an alternative empirical approach); this method has the advantage of being simple, model independent (i.e., it does not depend on the assumption of any particular set of galaxy spectral templates), and easy to use in determining approximate redshifts of $z\\la 4$ galaxies.   We determine the empirical analytic relations for color redshifts in \\S 2. We compare our estimated redshifts for the HDF galaxies with those obtained by the template-fitting method in \\S 3. We describe our estimated redshift catalog of HDF galaxies to $z\\la 4$ in \\S 4 (the Web site address of the catalog is given). We investigate the redshift clustering of HDF galaxies in \\S 5, and summarize our results in \\S 6. ",
        "conclusions": "Using HDF photometric and spectroscopic data, we have determined a set of simple analytic formulae that yield estimated galaxy redshifts to $z\\la 4$ in terms of linear combinations of three measured colors, $U-B$, $B-V$, and $V-I$ (Eqs.(\\ref{eq:Lz3,1})-(\\ref{eq:Hz2,2})). The derived analytic formulae in five color ranges exhibit small dispersions between the estimated and spectroscopic redshifts. For $z\\la 2$ galaxies, the redshift dispersion ranges from $\\sigma_z=0.034$ to $\\sigma_z=0.097$ for different color ranges. For $z\\ga 2$ galaxies, we find $\\sigma_z=0.14$ and $\\sigma_z=0.36$ for two color ranges which typically represent $z\\ga 3$ and $z\\la 3 $ galaxies respectively. These color-redshift relations apply to about 90% in the sample.  The smallest dispersion between the color and the spectroscopic redshifts, $\\sigma_z=0.034$, occurs for the $z\\la 2$ galaxies satisfying $(U-B)< (B-V) -0.1$; 28 galaxies with measured redshifts are used in deriving the relation for the estimated redshift, with only 4 free parameters (the coefficients in Eq.(\\ref{eq:Lz3,1})). There are 230 HDF galaxies with $I< 27$ and measured UBVI magnitudes that belong to this color range; we investigate the large-scale redshift distribution of these galaxies and find evidence for peaks in the redshift distribution that suggest large-scale clustering to at least $z\\sim 1$. These results are consistent with those of Cohen et. al. (1996) using observed spectroscopic redshifts of a smaller number of galaxies.  We have applied our color redshift formulae to the entire HDF photometric catalog and find that the derived redshifts are consistent with those obtained from spectral template-fitting techniques. The analytic relations, by design, yield lower dispersion than the template-fitting method. The color-redshift relations have the advantage of being simple, model independent, and easy to use. They can be further improved with additional data. These analytic color-redshift estimators are useful in providing empirical estimates of galaxy redshifts to $z\\la 4$ using multiband photometry.  Our Estimated Redshift Catalog of HDF Galaxies, based on our color redshift formulae for all 848 HDF galaxies with $I<27$ and measured $UBVI$ fluxes, is available by anonymous ftp in the elt/:HDF subdirectory of astro.princeton.edu.  Note that our color-redshift relations (Eqs.(\\ref{eq:Lz3,1})-(\\ref{eq:Hz2,2})) are derived using AB magnitudes and for the HDF filters. For application to other photometric catalogs, the appropriate spectroscopic training set should be used; when such a training set is not available, Eqs.(\\ref{eq:Lz3,1})-(\\ref{eq:Hz2,2}) may provide useful estimates after appropriate photometric transformation has been performed between the different filter systems. Also note that these color-redshift relations should not be applied to galaxies which lie close to the boundaries of the color ranges.  Finally, we note that our color-redshift relations are limited by the absence of measured spectroscopic redshifts for galaxies in the range of $1.4 \\le z \\le 2.2$ (see Fig.1 and Fig.2). It is very important to obtain spectroscopic redshifts in this range, because it will not only enable better calibration of photometric redshifts, it will also help us understand the nature of galaxies in the intermediate redshift range."
    },
    "9804/astro-ph9804148_arXiv.txt": {
        "abstract": "Catalysis of nucleon decay in white dwarfs is used to constrain the abundance of magnetic monopoles arising from Grand Unified Theories. Recent discoveries of the dimmest white dwarf ever observed, WD 1136-286 with $L = 10^{-4.94} L_{\\odot}$, place limits on the monopole flux. An abundance of monopoles greater than the new bound would heat this star to a luminosity higher than what is observed. The new bound is $(F/$cm $^{-2}$ s$^{-1}$ sr$^{-1}$) $(\\sigma \\upsilon/10^{-28} {\\rm cm}^2) < 1.3 \\times 10^{-20} (\\upsilon/ 10^{-3}c)^2$, where $\\upsilon $ is the monopole velocity. The limit is improved by including the monopoles captured by the main-sequence progenitor of the white dwarf: $(F/$cm $^{-2}$ s$^{-1}$ sr$^{-1}$ ) $(\\sigma \\upsilon /10^{-28} {\\rm cm}^2) < 3.5(26) \\times 10^{-21}$ for $10^{17}$ ($10^{16}$) GeV monopoles.  We also note that the dependence on monopole mass of flux bounds due to catalysis in neutron stars with main sequence accretion has previously been calculated incorrectly (previously the bound has been stated as $F (\\sigma \\upsilon/10^{-28} {\\rm cm}^2) <  10^{-28} $ cm $^{-2}$ s$^{-1}$ sr$^{-1}$).  We show that the correct bounds are somewhat weaker for monopole mass other than $10^{17}$ GeV. ",
        "introduction": "The question of whether or not magnetic monopoles exist has intrigued theorists and experimentalists for a long time \\cite{[1]}. In 1974, t'Hooft \\cite{[2]} and Polyakov \\cite{[3]} independently showed that magnetic monopoles always appear as stable topological entities in any Grand Unified theory (GUT) that breaks down to electromagnetism. Hence, if Grand Unified theories are shown to be correct, monopoles of mass in the range $10^{15}$ - $10^{19}$ GeV should exist. Rubakov \\cite{[4]} and Callan \\cite{[5]} calculated that these monopoles catalyze nucleon decay with a cross section characteristic of strong interactions, $\\sigma \\upsilon \\approx 10^{-28}$ cm$^2$.  The abundance of these monopoles is an open question. The Kibble mechanism predicts roughly one monopole per horizon volume at the time of the Grand Unified phase transition. However, this estimate provides a severe overabundance of the number of monopoles: monopoles overclose the Universe by many orders of magnitude. Instead  an inflationary epoch \\cite{[6]} may reduce their density in the Universe. Then the present abundance is difficult to estimate. A clue for experimentalists about what monopole flux to expect can be provided by astrophysics.  The Parker bound \\cite{[7]} on the flux of monopoles was obtained by requiring survival of $\\mu$G magnetic fields observed in our Galaxy and gave $F \\leq 10^{-16}$ cm$^{-2}$ sr$^{-1}$ sec$^{-1}$.  Subsequent improvements on this work include consideration of the monopole velocities \\cite{[8]} due to acceleration by the galactic magnetic field. Another improvement is the extended Parker bound, which required survival of a smaller seed magnetic field in the early period of the Galaxy \\cite{[9]}: $F \\leq 1.2 \\times 10^{-16} (\\frac {m}{10^{17}GeV})$ cm$^{-2}$ s$^{-1}$ sr$^{-1}$.  Another class of methods for determination of the monopole flux is based on the hypothesis that GUT monopoles give rise to the catalysis of nucleon decay. The basic idea is that monopoles traveling through the Galaxy lose enough energy to be captured in an object (e.q. white dwarfs, neutron stars, etc.) where they subsequently catalyze nucleon decay. The energy produced by the nucleon decay heats up the object and results in a flux of photons from the surface of the object. One can then compare this predicted luminosity with what is actually observed. One must ensure that the monopoles would not make the object brighter than what is seen. The coolest star (or other object) seen provides the tightest limit on the monopole flux. If there were more monopoles than allowed by the bound, then the dimmest star observed could not exist.  Several authors have carried out this kind of analysis in neutron stars \\cite{[10]}, nearby pulsar and white dwarfs. The strongest bound was obtained from consideration of the catalysis process in PSR 1929+10, an old pulsar \\cite{[11]}. From this pulsar, the bound on the product of monopole flux times cross section for catalysis is $(F/$cm$^{-2}$sr$^{-1}$sec$^{-1})(\\sigma\\upsilon / 10^{-28}$ cm$^{2}) \\leq 7 \\times 10^{-22}$. If one includes the monopoles captured by the main sequence progenitor of the white dwarf, this bound becomes even tighter \\cite{[12]}, \\hfill\\break $(F/$cm$^{-2}$ sr$^{-1}$ sec$^{-1})(\\sigma\\upsilon / 10^{-28}$ cm$^2$ $) \\leq 10^{-28}$. The consideration of monopole dynamics inside superconducting neutron-star cores leads to a bound $5 \\times 10^{-24} \\tau_{10}^{-2}$ cm$^{-2}$ sr$^{-1}$ s$^{-1}$ \\cite{[13]}, where $\\tau_{10}$ is the age (in $10^{10}$ years) of the pulsar's present magnetic field.  As neutron stars are the densest astrophysical objects observed, they give rise to the tightest catalysis bounds. However, there is a certain amount of uncertainty due to the fact that the interiors of neutron stars are not well understood. For example, neutron stars can have very large magnetic fields $\\sim 10^{12}$G of unknown topology, and the motion of magnetic monopoles inside the neutron star would undoubtedly be affected by these magnetic fields.  In addition neutron star interiors may contain pion condensates, again with uncertain effects on the monopoles. Because of the uncertainties with neutron star interiors, we turn to the next densest astrophysical objects in the Universe, white dwarfs. These stellar remnants are far better understood. The flux limits obtained from consideration of the catalysis process in white dwarfs are therefore important. Previously Freese \\cite{[14]} considered monopole catalyzed nucleon decay in white dwarfs. By comparing with the lowest luminosity white dwarf that had been seen at that time, she obtained a limit  \\begin{equation} (F/{\\rm cm}^{-2}{\\rm s}^{-1}{\\rm sr}^{-1}) (\\sigma\\upsilon /10^{-28}{\\rm cm}^2) \\leq 2\\times 10^{-18}. \\end{equation}  The present work is motivated by new observational data of cool white dwarfs \\cite{[15]}. In particular, Bergeron, Ruiz, and Leggett found a white dwarf 1136-286 (ESO 439-26) with luminosity $10^{-4.94}L_{\\odot}$; this is the dimmest white dwarf observed to date.  We use the measured luminosities of old white dwarfs to constrain the radiation due to monopole-catalyzed nucleon decay and thus to obtain an upper limit to the average flux of monopoles in the Galaxy. Since a white dwarf with luminosity $10^{-4.94}L_{\\odot}$ is observed today, we know that the monopole-induced contribution to the white dwarf luminosity cannot exceed this value. These new data improve the limit on the monopole abundance due to catalysis in white dwarfs \\cite{[14]} by roughly two orders of magnitude. Of course, as dimmer white dwarfs are found, the bound will continue to get more restrictive.  A monopole flux saturating this bound would keep the white dwarfs at luminosities at least this great and would lead to the prediction that no cooler white dwarfs will be found. As we will discuss, if it were indeed true that monopoles are keeping dwarfs hot, one would expect a different dependence of white dwarf luminosity on mass than expected in the standard model.  We shall explicitly indicate the dependence of our results on various parameters. We will parametrize the properties of the white dwarf in terms of typical values from observations: for the mass, $M = M_{0.6} 0.6 M_{\\odot}$, for the radius $R = R_{9} 9\\times 10^8$cm, and for the average density $\\bar \\rho = 4\\times 10^5$ g cm$^{-3} M_{0.6} R_{9}^{-3}$. The central density is about an order of magnitude higher, $\\rho_{c} = 3 \\times 10^6$ g cm$^{-3} M_{0.6} R_{9}^{-3}$. Rubakov \\cite{[4]} estimated the product of cross section for catalysis and relative velocity $\\upsilon$ of the monopole and nucleon to be constant: $\\sigma \\upsilon = \\sigma_{0} = 10^{-28} ({\\sigma\\upsilon}_{-28})$ cm$^2$. (Throughout, we take $\\hbar  = k_{B} = c = 1$.) For the thermal nucleon velocities expected inside a carbon and oxygen white dwarfs, $\\upsilon \\approx 10^{-3}c$, suppression effects may reduce the cross section by a factor of $10^{-2} s_{-2}$ \\cite{[16]}, and so we include this factor. In white dwarfs made of helium the suppression effects would be less effective $(s_{-2} = 10)$, and all the monopole flux bounds would be an order of magnitude stronger. ",
        "conclusions": "Figure 1 shows a plot of several monopole bounds: the Parker bound, the extended Parker bound, neutron star bounds, and the new white dwarf bound with and without main sequence capture.  In the plots we have used the Wood cooling curves to be conservative. We have found that consideration of newly observed white dwarf 1136-286 with luminosity $10^{-4.94} L_\\odot$ and with new calculations of white dwarf cooling curves leads to a bound on the monopole flux that is two orders of magnitude lower than previous bounds due to catalysis in white dwarfs.  The new bound is $F (\\sigma \\upsilon/10^{-28} {\\rm cm}^2) < 1.3 (1.9) \\times 10^{-20} (v/10^{-3}c)^2$ cm $^{-2}$ s$^{-1}$ sr$^{-1}$ for the Segretain \\cite{[19]} (Wood \\cite{[20]}) cooling curves respectively, where $\\upsilon $ is the monopole velocity. The limit is improved by including the monopoles captured by the main-sequence progenitor of the white dwarf: $F (\\sigma \\upsilon /10^{-28} {\\rm cm}^2) < 3.5(26) \\times 10^{-21}$ cm$^{-2}$ s$^{-1}$ sr$^{-1}$ for $10^{17}$ ($10^{16}$) GeV monopoles with $g=g_D$. Flux bounds for other monopole masses and parameters are given in Table 1. If cooler white dwarfs are discovered, a stricter bound on the monopole flux will result.  We also showed that the dependence on monopole mass of flux bounds due to catalysis in neutron stars with main sequence accretion has previously been calculated incorrectly.  Previously the bound due to catalysis in PSR 1929+10 with main sequence accretion has been stated as \\cite{[11]} $F (\\sigma \\upsilon/10^{-28} {\\rm cm}^2) <  10^{-28} $ cm $^{-2}$ s$^{-1}$ sr$^{-1}$. Instead, as can be seen in Table 1 and Figure 1, the correct bounds are somewhat weaker for monopole mass other than $10^{17}$ GeV.  {\\bf Figure Caption} Bounds on the monopole flux as a function of monopole mass. The Parker bound \\cite{[7]} due to survival of the galactic magnetic field is plotted, as is the extended Parker bound \\cite{[9]} due to survival of the magnetic field early in the history of the Galaxy. Mass density limits ($\\Omega h^2 <1$) are plotted for a uniform density of monopoles in the universe. The bounds due to catalysis in white dwarf WD1136-286 as discussed in this paper are plotted; the plots assume the cooling curves of Wood \\cite{[20]}, and are very similar to those obtained using cooling curves of Segretain ${\\it et al.}$ In addition, the bounds from this white dwarf with main sequence accretion (WD/MS) are plotted for $g=g_D$ (solid line) and for $g=2 g_D$ (dotted line).  The bounds due to calaysis in neutron star PSR 1929+10 are plotted, as are bounds due to this neutron star with main sequence accretion.  Again the solid line is for $g=g_D$ and the dotted line is for $g=2g_D$.  Note that the neutron star bounds with main sequence accretion have dependence on the monopole mass.  \\vfill\\eject \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|c|}\\hline &&&&&&$WD {\\rm w}/MS$&$NS {\\rm w}/MS$\\\\ \\cline{7-8} $M_m$ (GeV)&$\\beta$&${b_{\\rm crit}\\over R}$&$g/g_D$ &$N_{MS}/10^{38}F$&$N_{WD}/10^{38}F$&$F(\\sigma\\upsilon)_{-28}$&$ F(\\sigma\\upsilon)_{-28}$ \\\\[5pt] \\hline \\lower1.5ex\\hbox{$10^{15}$}&\\lower1.5ex\\hbox{$10^{-2}$} &{0.4}&{1}&2.5 &\\lower1.5ex\\hbox{0.17} &$8.2\\times 10^{-20}$& $6.2\\times 10^{-27}$ \\\\ \\cline{3-5}\\cline{7-8} &&0.56&{2}&4.9 && $4.3\\times 10^{-20}$& $3.2\\times 10^{-27}$ \\\\\\hline \\lower1.5ex\\hbox{$10^{16}$}&\\lower1.5ex\\hbox{$3\\times 10^{-3}$}&0.48&1&7.4&\\lower1.5ex\\hbox{1.8}&$2.4\\times 10^{-20}$& $2.1\\times 10^{-27}$ \\\\ \\cline{3-5}\\cline{7-8} &&0.62&2&12.3 && $1.6\\times 10^{-20}$& $1.3\\times 10^{-27}$ \\\\\\hline \\lower1.5ex\\hbox{$10^{17}$}&\\lower1.5ex\\hbox{$10^{-3}$}&0.54&1&52&\\lower1.5ex\\hbox{17}& $3.2\\times 10^{-21}$& $3.0\\times 10^{-28}$ \\\\ \\cline{3-5}\\cline{7-8} &&0.68&2&82 &&$2.2 \\times 10^{-21}$& $1.9\\times 10^{-28}$ \\\\ \\hline \\lower1.5ex\\hbox{$10^{18}$}&\\lower1.5ex\\hbox{$10^{-3}$}&---&1&---& \\lower1.5ex\\hbox{17 }&---&---\\\\ \\cline{3-5}\\cline{7-8}&&0.24&2&10 &&$8.1 \\times 10^{-21}$& $1.6\\times 10^{-27}$ \\\\\\hline \\end{tabular} \\end{center} \\bigskip \\begin{center} \\parbox{6in}{\\small Table 1: Bounds on the flux $F$ of magnetic monopoles in cm$^{-2}$s$^{-1}$sr$^{-1}$. Monopoles captured by white dwarfs (WD) or neutron stars (NS) and their main sequence (MS) progenitors are included. The white dwarf cooling time is taken to be $\\tau = 9.63$ Gyr.  $M_m$ is the monopole mass in GeV, $\\beta$ is the monopole velocity, ${b_{\\rm crit} \\over R}$ is the ratio of the critical impact parameter for a monopole in units of the radius of the main sequence star, and the monopole charge is $g= 69 e (g/g_D)$ in units of the Dirac charge $g_D$.  The number of monopoles captured by the MS progenitor and by the white dwarf are $N_{MS}$ and $N_{WD}$ respectively.  The second to last column is the flux bound due to catalysis in WD 1136-286 (with MS monopoles included).  The last column is the (corrected) flux bound due to catalysis in neutron star PSR 1929+10 (with MS monopoles included).} \\end{center} \\vspace{10mm} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|}\\hline $M_m$ (GeV)&$\\beta$&${b_{\\rm crit}\\over R}$&$g/g_D$&$N_{MS}/10^{38}F$&$N_{WD}/10^{38}F$&$F(\\sigma\\upsilon)_{-28}/ {\\rm cm}^{-2} {\\rm s}^{-1} {\\rm sr}^{-1}$ \\\\[5pt] \\hline \\lower1.5ex\\hbox{$10^{15}$}&\\lower1.5ex\\hbox{$10^{-2}$} &{0.4}&{1}&2.5 &\\lower1.5ex\\hbox{0.11} &$8.4\\times 10^{-20}$  \\\\ \\cline{3-5}\\cline{7-7} &&0.56&2&4.9 && $4.4\\times 10^{-20}$ \\\\\\hline \\lower1.5ex\\hbox{$10^{16}$}&\\lower1.5ex\\hbox{$3\\times 10^{-3}$}&0.48&1&7.4 &\\lower1.5ex\\hbox{1.2}&$2.6\\times 10^{-20}$ \\\\ \\cline{3-5}\\cline{7-7} &&0.62&2&12.3 && $1.6\\times 10^{-20}$ \\\\\\hline \\lower1.5ex\\hbox{$10^{17}$}&\\lower1.5ex\\hbox{$10^{-3}$}&0.54&1&52 &\\lower1.5ex\\hbox{11 }& $3.5\\times 10^{-21}$ \\\\ \\cline{3-5}\\cline{7-7} &&0.68&2&82 &&$2.4 \\times 10^{-21}$  \\\\ \\hline \\lower1.5ex\\hbox{$10^{18}$}&\\lower1.5ex\\hbox{$10^{-3}$}&---&1&---& \\lower1.5ex\\hbox{11 }&--- \\\\ \\cline{3-5}\\cline{7-7}&&0.24&2&10 &&$1.0 \\times 10^{-20}$  \\\\\\hline \\end{tabular} \\end{center} \\bigskip \\begin{center} {\\small Table 2: Same as table 1 for white dwarfs, but for cooling time $\\tau = 6.47$ Gyr.} \\end{center}   {\\bf Acknowledgements:} We thank J. Allyn Smith, D. Graff, G. Laughlin, G. Tarle, and V.D. Ivanov for helpful conversations. We acknowledge support from the DOE at the University of Michigan."
    },
    "9804/astro-ph9804238_arXiv.txt": {
        "abstract": "s{ We use numerical simulations of ray tracing through N-body simulations to investigate weak lensing by large-scale structure. These are needed for testing the analytic predictions of two-point correlators, to set error estimates on them and to investigate nonlinear gravitational effects in the weak lensing maps. On scales larger than 1 degree gaussian statistics suffice and can be used to estimate the sampling, noise and aliasing errors on the measured power spectrum. For this case we describe a minimum variance inversion procedure from the 2-d to 3-d power spectrum and discuss a sparse sampling strategy which optimizes the signal to noise on the power spectrum. On degree scales and smaller the shear and convergence statistics lie in the nonlinear regime and have a non-gaussian distribution. For this regime ray tracing simulations are useful to provide reliable error estimates and calibration of the measurements. We show how the skewness and kurtosis can in principle be used to probe the mean density in the universe, but are sensitive to sampling errors and require large observed areas. The probability distribution function is likely to be more useful as a tool to investigate nonlinear effects. In particular, it shows striking differences between models with different values of the mean density $\\Omega_m$. } ",
        "introduction": "Weak lensing by large-scale structure (LSS) shears the images of distant galaxies. The first calculations of weak lensing by LSS (Blandford et al. 1991; Miralda-Escude 1991; Kaiser 1992), based on the pioneering work of Gunn (1967), showed that lensing would induce coherent ellipticities of order 1$\\%$ over regions of order one degree on the sky. Recently several authors have extended this work to probe semi-analytically the possibility of measuring  the mass power spectrum and cosmological parameters from the second and third moments of the induced ellipticity or convergence (e.g. Bernardeau et al. 1997; Jain and Seljak 1997; Schneider et al. 1997).  The analytical work cited above suggested that nonlinear evolution of the density perturbations that provide the lensing effect can significantly alter the predicted signal. It enhances the power spectrum on scales below one degree and makes the probability distribution function (pdf) of the ellipticity and convergence non-Gaussian. We have carried out numerical simulations of ray tracing through N-body simulation data to compute the fully nonlinear moments and pdf. Details of the method and results are presented in a forthcoming paper; here we summarize the method and present some highlights of the results in Figures 1-5. We also discuss reconstruction of the dark matter power spectrum and error estimation using the gaussian approximation. ",
        "conclusions": ""
    },
    "9804/astro-ph9804242_arXiv.txt": {
        "abstract": "The spatial distribution of the Ly$\\alpha$ forest is studied using new HST data for the quasar pair Q~1026--0045 A and B at $z_{\\rm em}$~=~1.438 and 1.520 respectively. The angular separation is 36~arcsec and corresponds to transverse linear separations between lines of sight of $\\sim$300$h^{-1}_{50}$~kpc ($q_{\\rm o}$~=~0.5) over the redshift range  0.833~$<$~$z$~$<$~1.438. From the observed numbers of coincident and anti-coincident Ly$\\alpha$ absorption lines, we conclude that, at this redshift, the Ly$\\alpha$ structures have typical dimensions of $\\sim$500$h^{-1}_{50}$~kpc, larger than the mean separation of the two lines of sight. The velocity difference, $\\Delta V$, between coincident lines is surprisingly small (4 and 8 pairs with $\\Delta V$~$<$~50 and 200~km~s$^{-1}$ respectively). \\\\ Metal line systems are present at $z_{\\rm abs}$~=~1.2651 and 1.2969 in A, $z_{\\rm abs}$~=~0.6320, 0.7090, 1.2651 and 1.4844 in B. In addition we tentatively identify a weak Mg~{\\sc ii} system at $z_{\\rm abs}$~=~0.11 in B. It is remarkable that the $z_{\\rm abs}$~=~1.2651 system is common to both lines of sight. The system at $z_{\\rm abs}$~=~1.4844 has strong O~{\\sc vi} absorption.\\\\ There is a metal-poor associated system at $z_{\\rm abs}$~=~1.4420 along the line of sight to A with complex velocity profile. We detect a strong Ly$\\alpha$ absorption along the line of sight to B redshifted by only 300~km~s$^{-1}$ relatively to the associated system. It is tempting to interpret this as the presence of a disk of radius larger than 300$h^{-1}_{50}$~kpc surrounding quasar A. ",
        "introduction": "\\label{intr} One way to probe the transverse extension of the gaseous structures giving rise to the Ly$\\alpha$ forest seen in the spectrum of quasars is to observe multiple lines of sight to quasars with small angular separations on the sky and search the spectra for absorptions coincident in redshift.   This technique originated with a suggestion by Oort (1981) to test the possibility that the Ly$\\alpha$ forest clouds originate in large pancake structures. The first discoveries of common and associated absorption using pairs of distinct quasars (with separations $\\sim$~1~arcmin) were made by Shaver et al. (1982) and Shaver \\& Robertson (1983). These already indicated the possible existence of very large absorber sizes (hundreds of kpc), even for the Ly$\\alpha$ clouds. At about the same time Sargent et al. (1982) found no detectable tendency for Ly$\\alpha$ lines to correlate in QSO pairs separated by a few arcmin. Spectra of pairs of gravitational lens images revealed common absorptions on smaller scales (Weyman \\& Foltz 1983, Foltz et al. 1984). The idea that Ly$\\alpha$ clouds might have large sizes remained controversial untill the analysis by Smette et al. (1992), later confirmed by Dinshaw et al. (1994), Bechtold et al. (1994), Crotts et al. (1994), Bechtold \\& Yee (1994), Smette et al. (1995), D'Odorico et al. (1998).  Recently, Dinshaw et al. (1995) derived a radius of 330$h^{-1}_{50}$ kpc at $z$~$\\sim$~0.7 for spherical clouds from observation of Q0107--0232 and Q0107--0235 separated by 86~arcsec. Larger separations have been investigated by Crotts \\& Fang (1997) and Williger et al. (1997). Both studies conclude that the clouds should be correlated on scales larger than 500~kpc.\\par\\noindent Here we present observations of Q1026--005~A ($m_{\\rm r}$~=~18.4, $z_{\\rm em}$~=~1.438) and B ($m_{\\rm r}$~=~18.5, $z_{\\rm em}$~=~1.520), two distinct quasars separated on the sky by 36~arcsec or 300~$h^{-1}_{50}$~kpc ($q_{\\rm o}$~=~0.5) at $z$~$\\sim$~1. ",
        "conclusions": ""
    },
    "9804/astro-ph9804304_arXiv.txt": {
        "abstract": "How has the ``fluffy'' core of the Sgr dwarf galaxy survived multiple strong shocks from the tidal force of the Galactic halo and disc since the formation of the core a Hubble time ago?  A scenario that Sgr was deflected to its current orbit by the Magellanic Clouds after a rendezvous on the north Galactic pole $2-3$ Gyrs ago is examined.  It is shown that the conditions of the collision fix both the sense of circulation of Sgr and the LMC around the Galaxy and the slope of the Galactic rotation curve.  The model argues that the two orthogonal polar circles traced by a dozen or so Galactic halo dwarf galaxies and globular clusters (LMC-SMC-Magellanic Stream-Draco-Ursa Minor along $l \\approx 270^o$ and M54-Ter 7-Ter 8-Arp 2-NGC 2419-Pal 15 along $l \\approx 0^o$) are streams of tidal relics from two ancient galaxies which was captured on two intersecting polar rosette orbits by the Galaxy.  Our results favor the interpretation of microlensing towards the LMC being due to source or lens stars in tidal features of the Magellanic Clouds.  We discuss direct and indirect observations to test the collision scenario. ",
        "introduction": "The recently discovered dwarf galaxy at about $25$ kpc from the Sun in the direction of the Sagittarius constellation (Ibata, Gilmore \\& Irwin 1994) is the closest galaxy known to us.  It is traced by two long trailing/leading tails on the sky (together more than $8^o\\times 22^o$ in solid angle) with most of its stars still clustered around a low density luminous core (roughly $0.001 L_\\odot{\\rm pc}^{-3}$ with semi-axes $1:1:3$ kpc).  It is puzzling why this fluffy core of the dwarf galaxy has not been fully ``digested'' by the Galaxy, in the sense that stars have not fully dispersed out of the core despite the severe shocks at pericentric passage from the tidal force of the Galactic halo (about $10-100$ times stronger than that experienced by satellites in the outer halo, the Magellanic Clouds and the Fornax dwarf galaxy included) and shocks when crossing the disc of the Galaxy.  The best fit to Sgr's morphology, radial velocity (Ibata, Gilmore \\& Irwin 1995) and proper motion (Ibata, Wyse, Gilmore \\& Suntzeff 1997) yields an orbit with a pericenter-to-pericenter period of about $0.8$ Gyr and a peri and apo-center at about 10 and 50 kpc respectively (Vel\\'azquez \\& White 1995).  Simulations show that if a typical Galactic dwarf galaxy (such as Fornax) were replaced on Sgr's orbit, it would dissolve in no more than two peri-centric passages by the strong peri-centric tidal shock of the Galaxy near 10 kpc (Vel\\'azquez \\& White 1995; Johnston, Spergel \\& Hernquist 1995; Johnston, Hernquist \\& Bolte 1996; Edelsohn \\& Elmegreen 1997).  This apparently contradicts the observation that the dominant stellar population in the core is older than 10 Gyrs (Mateo et al. 1995, Fahlman et al. 1996), implying that Sgr has survived $10-20$ peri-centric tidal shocks of the Galaxy.  To circumvent this dilemma we need to abandon either or both of the following hidden assumptions: (i) the light distribution of Sgr traces its mass, and (ii) Sgr has always been on the same low-pericentric orbit in a rigid Galactic potential for the past 5 to 10 Gyrs.  Ibata, Wyse, Gilmore \\& Suntzeff (1997) postulate a dense dark halo of Sgr surrounding the luminous part to hold the system together; they require Sgr's mass density to be uniform inside about 3 kpc of its core with a value ($\\sim 0.03 M_\\odot{\\rm pc}^{-3}$) several times the mean Galactic halo density inside $10$ kpc ($0.013M_\\odot{\\rm pc}^{-3}$).  An inspection of Sgr's rosette-like orbit in relation to that of the Magellanic Clouds (MCs) offers a completely different line of thought.  They are on nearly orthogonal planes intersecting along the poles with their Galactocentric radii overlapping at about 50 kpc. So an encounter at the north or south pole some time in the past or future is quite inevitable.  A recent preprint by Ibata \\& Lewis (1998), shortly after the completion of the work reported in this {\\it Letter}, also remarked on a small chance of an interaction after noticing in their simulations a weak perturbation to Sgr's orbit when they turned on the moving gravitational field of the massive MCs. Unfortunately the effect was in the end neglected on grounds of low probability without thoroughly exploring the parameter space (of satellite velocities and the Galactic potential) and the important consequences of a rare strong interaction.  So same as Ibata et al. (1997) they were thus left with no alternative but to conclude a massive dark halo of Sgr to be the only explanation for Sgr's presence on a low-pericentric orbit after a Hubble time.  In this {\\it Letter} we examine the encounter scenario, as illustrated in Fig. 1, where Sgr has been pulled back from an originally high angular momentum/energy orbit to the present low angular momentum/energy orbit by the massive MCs.  A recent encounter would have the advantage to allow Sgr to spend most of its lifetime on a ``safe'' orbit with a pericenter (say, $20$ kpc) too high to be harassed by the sharply declining tidal force of the Galaxy (e.g., Kroupa 1997, Oh, Lin \\& Aarseth 1995); in a halo with an $r^{-2}$ density profile the pericentric shock would drop by a factor of $4$ from $10$ kpc to $20$ kpc.  \\onecolumn \\begin{figure} \\epsfysize=15cm \\centerline{\\epsfbox{sgrmc3d.ps}} \\caption{ A 3D view and $x-y$, $y-z$ projections of the orbits of the LMC/SMC (thick red/dashed green lines near the $x=0$ plane) and Sgr (thin blue curve near the $y=0$ plane); the Sun is be at $(x,y,z)=(-8,0,0)$ kpc. The three systems are integrated backward for $3$ Gyrs (with ellipses marking steps of $0.5$ Gyr) from the present epoch (with velocity vectors marked by arrows) inside the Galactic potential.\\label{sgrmc3d} } \\end{figure} \\twocolumn  Various interesting aspects of this scenario will be discussed at the end.  But the aim of this {\\it Letter} is to report an independent constraint on the rotation curve of the Galaxy as imposed purely by {\\it timing the collision}.  The essence is the following.  The random chance for the LMC and Sgr to meet each other is obviously low, about 1\\% for a 10 kpc closest approach in the past 3 Gyrs for a general set of Galactic potentials and initial conditions of the satellites.  So the same argument could be inverted: once we accept the deflection by the MCs as a plausible way out of Sgr's dilemma a stringent set of conditions on the potential of the halo and the proper motions of the satellites must follow. ",
        "conclusions": ""
    },
    "9804/astro-ph9804214_arXiv.txt": {
        "abstract": "A new approach to the study of the large-scale stellar cluster distribution in the Galaxy based on two-point correlation techniques is presented. The basic formalism for this method is outlined and its applications are then investigated  by the use of a simple   model of cluster distribution in the Galaxy. This provides an  estimate of  the potentials of the two-point correlation function for indicating clustering in the measured star positions, which can be related to the presence of star clusters in the observed field. This technique is then applied to several areas of the Two Micron Galactic Survey catalogue, from which information is obtained on the distribution of clusters  according to position in the Galaxy, as well as about age, density of stars, etc. ",
        "introduction": "Open  star clusters provide valuable information on the evolution of the Galaxy. In this paper two-point correlation techniques are used to analyse the distribution of open clusters in order to  gain an insight into the structure and evolution of the Galaxy.  Open-cluster distributions have been widely studied at optical wavelengths as a means of studying Galactic structure and evolution (see, for example, Lyng\\aa \\ 1987b; Janes \\& Phelps 1994). The Lyng\\aa\\ catalogue of open clusters (Lyng\\aa \\ 1987a) lists about 1200 clusters,  which represent nearly all the open clusters accessible in the  visible. Knowledge of the positions and ages of these clusters (a method of age determination for clusters is given by Carraro \\& Chiosi 1994) enables the scale length and scale height of the disc to be derived for both young and old clusters (for a review of old open clusters, see Friel 1995) and theories to be developed on their formation and destruction throughout the history of the Galaxy.  The limitations on these studies are imposed by the maximum distance at which open clusters can be detected. Most of the cataloged open clusters are in the solar neighbourhood,  and very few  have distances greater than 3 kpc (Payne-Gaposchkin 1979). Hence, information is obtained only for a small region of the Milky Way. The problem is caused by interstellar extinction.  An excellent tool for studying star clusters and star formation regions is to observe in the infrared (Wynn-Williams 1977), where the effect of the extinction is far less. To date, however, the infrared  has been little used in this field owing the absence of suitable databases.  The $K$ band is probably the best region of the spectrum for tracing the stellar distribution of the Galaxy. The $K$ radiation is a mass tracer in spiral galaxies because it follows the old stellar population (Rhoads 1995). Furthermore, the $K$ light is dominated by high-mass stars  in star formation regions, i.e. in open clusters, so it is especially appropriate in the search for young clusters, which are rich in massive stars.  As explained in more detail by Garz\\'on et al. (1993), the Two micron Galactic survey (TMGS) is a $K$-band survey of various regions along the Galactic equator between  $-5^o<l<35^o$, $|b|\\le 15^\\circ$ and $35^\\circ<l<180^\\circ$, $|b|\\le 5^\\circ$. The TMGS catalogue has a positional accuracy of about 4 arcsec in right ascension and 7 arcsec in declination. These errors have been estimated after cross-correlating the original TMGS  source positions with  {\\it Guide Star Catalogue} (GSC) counterparts. The larger error in declination comes from the orientation of the array with respect to the survey direction. Due to the dead spaces between detectors, the $K$ limiting magnitude for completeness has to be set conservatively within a range from 9 to 9.5 mag, although the detection limit magnitude of the survey is well in excess of 10 mag.  In this paper a new method is presented for automatically determining the level of clustering in catalogues,  the TMGS being used as an example. A set of  criteria are defined for an automatic search for correlations among stars by means of  the two-point correlation function and the two-point angular correlation function. A simple model which assumes  a regular distribution of clusters with a constant star density  is developed in Section 3. Predictions from this model are then compared with the TMGS and the discrepancies analysed. The causes of clustering are then discussed. The use of the tools described in Section 2 and their application to the TMGS catalogue are dealt with in Section 5, and the results for several regions of our Galaxy are given in Section 6. Finally,  a summary of the results is given, and suggestions are made for future developments of the methodology described here.  \\section {The two-point correlation function and the two-point angular correlation function}  Occasionally, when the scale length of a system is changed, certain aspects of the system remain invariable, as in the case for the distribution of matter in space. For example, there are mathematical methods for handling the spatial distribution of atoms in solids, gases and (particularly) liquids. Cosmologists face the same kind of mathematical problem when working with the distribution of galaxies and clusters of galaxies in the context of the large-scale structure of the Universe. They treat the Universe as a fluid whose `particles' are galaxies. Our aim in this paper is to develop the use of similar mathematical methods on an intermediate scale, i.e. in examining the distribution of the stars that make up our Galaxy.  Correlation functions describe how points are distributed (e.g. Peebles 1980; Borgani 1995). Suppose  a local density of objects $\\delta N/{\\delta V}=\\rho(\\vec{r})$ and an average density $\\langle\\rho\\rangle$ (hereafter, the symbol $\\langle\\cdots\\rangle$ indicates a local volume average).  Note that $\\rho $ corresponds to the exact distribution of objects, i.e. it is a Dirac delta function with  zero value where there are no objects, and $\\langle\\rho\\rangle$ is the local average density, i.e. the number of objects per unit volume, and provides no information concerning their distribution.  The two-point correlation function (TPCF) is defined as  \\begin{equation}\\xi (|\\vec{r}-\\vec{r}'|)=\\frac{\\langle\\rho(\\vec{r}) \\rho(\\vec{r}')\\rangle} {\\langle\\rho(\\vec{r})\\rangle^2} -1 .\\label{xi}\\end{equation}  The function $\\xi (r)$ expresses the excess over the random probability of finding objects at separation $r$. ($\\xi (r)=0$ means that the probability is totally random; $\\xi (r)>0$  that the probability is greater than random, i.e. that there is clustering; and   $\\xi (r)<0$ that the probability is less than random, i.e. that there is relative avoidance).  In the same way, the two-point correlation function can be defined for two  dimensions on the surface on to which the distribution is projected (the celestial sphere in the case considered here). This is called the two-point angular correlation function (TPACF) and is defined as  \\begin{equation}\\omega (|\\vec{\\theta}-\\vec{\\theta}'|)=\\frac{\\langle\\sigma (\\vec{\\theta}) \\sigma(\\vec{\\theta}')\\rangle} {\\langle\\sigma(\\vec{\\theta})\\rangle^2} -1 \\label{omega},\\end{equation} where $\\sigma $ is the surface density per unit solid angle.  Another mathematical technique for deciding whether a distribution is non-Poissonian is area tessellation, as was used by Bal\\'azs (1995) to test the grouping tendency of H$_\\alpha$-emission stars in the Orion molecular clouds without giving a quantitative measure of the departure from the Poissonian distribution. See also P\\'asztor et al. (1993), P\\'asztor \\& T\\'oth (1995) and references therein for other astronomical applications of spatial statistics.  \\subsection{Relationship between the TPCF and the TPACF for stars}  When applying the above definitions to stars in the Galaxy, the luminosity function and space density have to be taken into account.  By generalizing the result of the Limber (1953) equation for constant density, the relationship between the TPCF (which is non-zero for distances less than $\\Delta r$) and the TPACF (for small $\\theta$) for any density distribution is  \\[ \\omega _{\\rm t}(\\theta)\\approx \\frac{1}{\\langle\\sigma _{\\rm t}\\rangle^2} \\int _0^{\\infty }dr\\ r^4 \\langle\\rho \\rangle^2(r) \\int_{r-\\Delta r}^{r+\\Delta r}dr' \\]\\[ \\times \\int _{M_{\\rm min}(r)}^{M_{\\rm max}(r)} dM\\ \\phi(M)\\int_{M_{\\rm min}(r)}^{M_{\\rm max}(r)}dM'\\ \\phi(M') \\]\\begin{equation} \\times \\xi \\left(\\sqrt{ (r\\theta )^2+(r-r')^2};r,M,M'\\right) ,\\label{omegat2}\\end{equation} where $r$ is the distance along the line of sight, $M$ the absolute magnitude, $\\phi (M)$ the lu\\-mi\\-no\\-si\\-ty function and  \\begin{equation} \\langle\\sigma _{\\rm t}\\rangle=\\int _0^{\\infty }dr\\ r^2\\langle\\rho \\rangle(r)\\int _{M_{\\rm min}(r)} ^{M_{\\rm max}(r)} dM\\ \\phi(M) .\\label{sigmat2}\\end{equation} The minimum and maximum values of $M$ for a distance $r$ depend on the minimum and maximum values of the apparent magnitude and the extinction along the line of sight.  In this case, it is assumed that the absorption is not patchy, i.e. that it is independent of $\\theta $ for small angles. This is not exactly true but it will be  show in Section \\ref{.extincorr} that the effects are negligible.  The subscript `t' stands for `total', a projection of all distances and magnitudes, and $\\langle\\sigma _{\\rm t}\\rangle$ is the total two-dimensional density for all distances and magnitudes. In the literature, $\\langle\\sigma _{\\rm t}\\rangle$ is also called $A(m_{\\rm min},m_{\\rm max},l,b)$ and represents the star counts in the magnitude range (Bahcall 1986).  This expression enables the TPACF to be found once the three-dimensional distribution of the stars is known and forms the basis of this article, in which we create a model distribution of the stars and compare the results obtained with those observationally in order to investigate the distribution of clustering in the structure of our Galaxy.     In general, the TPACF cannot be inverted to give the TPCF due to the multiplicity of possible solutions and to the lack of precise knowledge of certain parameters. However, there are certain cases in which the equation can be inverted and TPCF obtained from TPACF (Fall \\& Tremaine 1977). A trivial example where inversion is possible is that of a Poissonian three-dimensional distribution, which implies a Poissonian projected distribution and vice versa, i.e. $\\xi =0$, $\\omega =0$ on all scales. Another example is when  $\\langle\\rho\\rangle(r)$ is a constant independent of $r$.  \\subsection{ Definition of new variables}  In order to simplify the comparison of the level of clustering for different regions of the sky, two new variables will be introduced.  $\\theta _{\\rm max}$ is defined as the first zero of $\\omega _{\\rm t}(\\theta)$. In this article (see for example Fig. \\ref{Fig:a052}), $\\omega _{\\rm t}$ is positive up to a separation $\\theta _{\\rm max}$. For values greater than $\\theta _{\\rm max}$ this value is small and oscillates about zero, as there is no correlation among stars separated by large angular distances.  Another definition, corresponding to the integration of $\\omega _{\\rm t}$ up to the limit $\\theta =\\theta _{\\rm max}$ ($\\theta >\\theta _{\\rm max}$ would give a null contribution to the integral), is  \\begin{equation} C_2\\equiv \\frac{\\int _0 ^{\\theta _{\\rm max}}d\\theta \\ \\theta \\omega _{\\rm t}(\\theta ) }{\\theta _{\\rm max}^2} ,\\label{C2}\\end{equation} which means the excess (when $C _2$ is positive) or deficit (when $C _2$ is negative) of the relative number of objects with respect to a Poissonian distribution in a circle centred on an arbitrary star on the celestial sphere, within the observed solid angle and with angular radius $\\theta _{\\rm max}$. The relative  correlation within the angular scale $\\theta _{\\rm max}$ is therefore measured. We call $C_2$ the `accumulation parameter' (N.B. there are also other definitions in the literature of the TPACF integral, e.g. Wiedemann \\& Atmanspacher 1990).  The variable $C_2$ has a clear meaning associated with projected clustering and is also a useful number to measure. Since it sums several values of $\\omega $ for different angles, it condenses the information of interest into a single number that can be compared for different samples of stars and give the degree of clustering. This parameter is a mathematical expression of the degree of clustering seen in fields of stars. The idea that we wish to stress here is that all mathematical developments  described in this paper are designed to put the intuitive idea of clustering to a reliable test. These calculations are necessary for a quantitative, as opposed to a merely qualitative, description of clustering.  Applying the expression (\\ref{omegat2}) of $\\omega _{\\rm t}$ in $C_2$, we get  \\[ C_2=\\frac{1}{\\langle\\sigma _{\\rm t}\\rangle^2\\theta _{\\rm max}^2} \\int _0^{\\infty }dr\\ r^2\\langle\\rho (r)\\rangle^2 \\int _{M_{\\rm min}(r)}^{M_{\\rm max}(r)} dM\\ \\phi(M) \\]\\begin{equation} \\times \\int _{M_{\\rm min}(r)}^{M_{\\rm max}(r)} dM'\\ \\phi(M') \\int _0^{\\theta _{\\rm max}r} dy\\ y \\Xi (y;r,M,M') \\label{C2a},\\end{equation} where $\\Xi $, an integrated TPCF, is  \\[ \\Xi(y;r,M,M') \\]\\begin{equation}= \\int_{r-\\Delta r}^{r+\\Delta r}dr' \\xi\\left (\\sqrt{y^2+(r-r')^2};r,M,M'\\right) .\\label{Xi}\\end{equation}   \\subsection{ Further approximations}  In order to simplify the above calculations, it will be assumed that the distribution of stars does not depend on their luminosity, i.e. that $\\xi (y;r,M,M')=\\xi (y;r)$. This is not completely true as there is a small dependence on the distribution of stars in a cluster according their masses, and the luminosities are dependent on the masses. A complete calculation taking the luminosity function  into consideration would be of great value. However, the relationship between the TPCF and the luminosity function is uncertain, although the effects of this approximation for the detection of clusters are expected to be small.  With this approximation,  and from (\\ref{sigmat2}) and (\\ref{omegat2}),  \\begin{equation} \\langle\\sigma _{\\rm t}\\rangle=\\int _0^{\\infty}dr\\ \\langle N^*\\rangle(r) \\end{equation} and \\begin{equation} \\omega _{\\rm t}(\\theta)=\\frac{1}{\\langle\\sigma _{\\rm t}\\rangle^2} \\int _0^{\\infty}dr\\ \\langle N^*\\rangle^2(r) \\Xi (r\\theta ;r) ,\\label{omegat*}\\end{equation} where $\\langle N^*(r)\\rangle$ is the number of stars observed per unit solid angle at a distance $r$:  \\begin{equation} \\langle N ^*\\rangle(r)=r^2\\langle\\rho \\rangle(r) \\int _{M_{\\rm min}(r)}^{M_{\\rm max}(r)} dM\\ \\phi(M) .\\label{N*}\\end{equation}  The variable $\\omega _{\\rm t}$ can also be expressed as  \\begin{equation} \\omega _{\\rm t}(\\theta)=\\overline{\\Xi (\\overline{r}\\theta )} \\label{omegaavXi}\\end{equation} where the averages $\\overline{r}$ and $\\overline{\\Xi }$ are such that match (\\ref{omegat*}).  Also, from (\\ref{C2a}),  \\begin{equation} C_2=\\frac{1}{\\langle\\sigma _{\\rm t}\\rangle^2\\theta _{\\rm max}^2} \\int _0^{\\infty }dr\\ \\frac{\\langle N^* (r)\\rangle^2} {r^2} \\int _0^{\\theta _{\\rm max}r} dy\\ y \\Xi (y;r) .\\label{C2*}\\end{equation}  \\noindent This last equivalence is a way of averaging the function $\\xi $.  Hence, high values of $C_2$ indicate that there must be high projected clustering in the direction of the beam.    \\subsection{Patchiness of extinction} \\label{.extincorr}  It is clear that extinction can distort the observed counts, the amount of the distortion being a matter of controversy. It is  generally accepted that in the optical wavelengths this influence is very severe, particularly in regions near to or in the Galactic plane in the inner Galaxy, where the strong and patchily distributed obscuration makes it difficult to penetrate deep into the Galaxy. The amount of extinction decreases substantially with  increasing wavelength. Maihara et al. (1978) quoted a value of 0.17 mag kpc$^{-1}$ as  typical for  extinction in the Galactic plane in the $K$ band, compared with 1.9 mag kpc$^{-1}$ for the $V$ band (Allen 1973).  This has two important consequences. First, the $K$ band is more effective at penetrating the interstellar dust. Secondly,  the observed stellar dis\\-tri\\-bu\\-tion more closely resembles the true distribution. For the second argument to be true it is necessary that the obscuration in the $K$ band should not only be smaller in amount than in the $V$ band, but also that its patchiness should be less important.  This rather uniform distribution of the interstellar extinction in $K$ can be inferred from the TMGS histograms in several cuts across the Galactic plane. Garz\\'on et al. (1993, their Fig. 8) compared the observed stellar distribution in the TMGS and the GSC in the $V$ band. It is noticeable how uniform the $K$ histograms are, particularly when  compared with those for the GSC. Except for small portions highly concentrated in the Galactic plane and more marked in the central regions, the shape of the high spatial resolution distribution curves of the TMGS does not exhibit the `noisy' pattern of the GSC plots, which is certainly due to the presence of strong and patchily distributed extinction.  Hammersley et al. (1994)   showed similar histograms for different areas which also have similar shapes. Moreover, a good fit to a classical ex\\-po\\-nen\\-tial disc can be seen in Fig. 3 of that paper; this would not be the case if the extinction were important and non-uniform.  This conclusion can also be reached from the contour maps of the bulge of the Galaxy of Dwek et al (1995), who showed the residuals of the DIRBE data after disc subtraction and extinction correction. Again, the general shape of the maps proves the basic uniformity of extinction distribution in the near infrared.    We now estimate these effects. From (\\ref{sigmat2}) with the change of variable $r=10^{(5+m_{\\rm max}-M_{\\rm max})/5}$ and  \\begin{equation} \\Phi (M_{\\rm max})=\\int _{-\\infty} ^{M_{\\rm max}} dM\\ \\phi(M) \\approx \\int _{M_{\\rm min}}^{M_{\\rm max}} dM\\ \\phi(M) ,\\end{equation} the local cumulative counts $\\sigma _{\\rm t}$ follow the  expression  \\[ \\sigma _{\\rm t}=\\langle\\sigma _{\\rm t}\\rangle _{local} =200(\\ln\\ 10)10^{3m_{\\rm max}/5}\\omega \\int_{-\\infty }^\\infty dM_{\\rm max} \\]\\begin{equation}\\times D\\left( 10^{(5+m_{\\rm max}-M_{\\rm max})/5}\\right) 10^{-3M_{\\rm max}/5} \\Phi (M_{\\rm max}), \\end{equation} ignoring the variation of extinction with the distance. If we take the density $D$ as  constant, then  \\begin{equation} \\sigma _{\\rm t}= N(m_{\\rm max})\\propto 10^{3m_{\\rm max}/5} .\\end{equation} Taking $D$ as  constant is sufficient for  estimating the  the order of magnitude of the patchiness due to extinction. In any case, the above proportionality is followed in the observed cumulative counts but with a constant value of between 1 and 2 instead of $3/5$ in the exponent.  An excess of extinction, $\\Delta a(\\theta)$, due for example to a cloud at an angular distance $|\\vec{\\theta } -\\vec{\\theta _0}|$ with respect to a given point $\\theta _0$, will cause a reduction in the apparent flux of a fraction, $f$, of stars (behind the cloud), thereby creating the same effect as a  reduction in the maximum apparent magnitudes of these stars by $\\Delta a(\\theta)$ mag, or, if $\\Delta a(\\theta)$ is relatively small, a  reduction in $m_{\\rm max}$ by $f\\Delta a(\\theta)$ mag for all stars. Hence,  \\begin{equation} \\sigma _{\\rm t}(\\theta)\\sim \\sigma _{\\rm t} (\\theta _0)10^{-3(a(\\theta )-a(\\theta _0))f/5} \\label{sigmada} .\\end{equation}  If it is assumed that the observed flux fluctuations, $\\Delta F$, are due mainly to extinction variations, with the small-fluctuation approximation, then  both are related by  \\begin{equation} \\Delta a=-2.5\\log \\left( 1-\\frac{1}{f}\\frac{\\Delta F}{F} \\right)\\approx \\frac{5\\log _{10}e}{2f}\\frac{\\Delta F}{F} \\end{equation} (the factor $f$ appears again here for the same  reasons as above). So, from  equation (\\ref{sigmada}), using the small-fluctuation approximation,  \\begin{equation} \\frac{\\sigma _{\\rm t}(\\theta )} {\\sigma _{\\rm t} (\\theta _0)} \\sim \\frac{3}{2} \\frac{F(\\theta)}{F(\\theta _0)} \\label{omegaext} .\\end{equation} This means that the angular correlation of star density is about 3/2 times the angular correlation of the flux.  Averaging the DIRBE $K$ flux (Boggess et al. 1992) fluctuations  from the maps with $2520''$ resolution over $\\mid b\\mid \\le 3^\\circ$ for constant-$l$ strips over the range $-35^\\circ < l< 35^\\circ $ (where the effects of extinction are most relevant), we get root mean squares of  \\begin{equation} 0.03< \\sigma (\\Delta F/F)_{l={\\rm const.}, \\mid b \\mid \\le 3^\\circ} < 0.23 ,\\end{equation} with an average of \\begin{equation} \\overline{\\sigma (\\Delta F/F) _{l={\\rm const.}, \\mid b \\mid \\le 3^\\circ}}= 0.10 .\\end{equation} The oscillations of flux fluctuations are not very high in the plane, their maximum being $2.3$ times the average.  From equation (\\ref{omegaext}), and taking into account that the root mean square is $\\sqrt{\\omega (0)}$,  \\begin{equation} \\overline{\\omega (0) _{l={\\rm const.}, \\mid b \\mid \\le 3^\\circ}} \\sim 0.015 \\end{equation} for regions of 2520$''$ in size. In the most unfavourable case, where the extinction is highest (multiplied by a factor of 2.3$^2$ because the maximum root mean square is $2.3$ times greater than the average), $\\omega (0) \\sim 0.08$.  Higher-resolution flux maps are not available in the  $K$-band for the whole sky so we cannot derive these numbers for smaller scales, but they are not expected to be much higher since average cloud size is of the order of degrees (rather higher than 42$'$) and the cloud distribution is fairly smooth. A fractal distribution would increase the contamination but this may apply only to very cold gas clouds (Pfenninger \\& Combes 1994) which are not the main cause of extinction in the  $K$ band.  We conclude that extinction in $K$ cannot be responsible for correlations $\\omega (0)$ greater than $\\sim 0.08$. This is just an estimate, but the order of magnitude should not be very different. As will be shown, the results when applied to the TMGS are above this value (see, for example, Fig. \\ref{Fig:a052}), and causes other than patchy extinction must explain this. ",
        "conclusions": "A technique is developed for  searching for clustering in stellar  surveys using  correlation functions. The mathematical tools are useful for any field of stars and can be applied to any survey, especially those at carried out at infrared wavelengths, which permit a study of the distribution of stars throughout almost the entire Galaxy. The DENIS (Epchtein 1997) or 2MASS (Skrutskie et al. 1997) surveys will be ideal for this technique as the increased numbers of stars will reduce the errors. It is even possible, with a large number of stars in the survey, to apply the technique for different ranges of apparent magnitude. Studying the clustering of stars at different apparent magnitudes is equivalent to do  studying in three dimensions ($l$, $b$ and the average distance $\\overline{r}$ which is associated with the treated range of magnitudes).  A simple model has been developed. This model could be improved by introducing a density dependence as a function of the distance from the centre of the cluster, perhaps a power-law dependence.  In this paper the method has been applied to the TMGS. Is has been shown that a simple model in which old open clusters trace the whole Galaxy with a density of clusters proportional to the density of stars agrees quite well with the data. An exception to the general agreement  are specific regions in the plane where the higher-than-expected clustering can be a\\-ttri\\-bu\\-ted  to star formation in the spiral arms. A second departure from the simple model  is the reduced $C_2$ in  the outer disc and in the bulge due to a lack of young  clusters.  In one of the regions with an excess, at $l=70^\\circ $ in the plane, the approximate limits for the cluster density and the density of stars inside the cluster are derived. These are, respectively, $5\\times 10^{-7}\\ {\\rm pc}^{-3} <{\\rm \\langle}n_{\\rm cl}{\\rm \\rangle}< 2 \\times 10^{-6}\\ {\\rm pc}^{-3}$ and $140\\ {\\rm pc}^{-3} < \\rho _{\\rm cl} < 700\\ {\\rm pc}^{-3}$. There is, however, a lower-than-expected  correlation  at $l=27^\\circ $, $b=0^\\circ $. There is believed to be a huge star formation  region in this direction and the lack of correlation could be due to the star formation region being far larger than the sample area.  As has been pointed out by Friel (1995), the oldest open clusters may be viewed from two perspectives with regard the formation of the Galaxy: a halo collapse or a continuous accretion and infall of material from the halo on to the Galactic disc. Either perspective is possible. The first should justify which were the original star formation regions that were the origin of the present old clusters in the outer disc and how they travelled there from their place of  origin. The second perspective needs  to test the infall of matter from the halo as well as the existence of star systems in the halo. Further improvements on these cluster searches and better numbers will give us a hint concerning these questions on the origin of the Galaxy. A better determination of $C_2$ in the bulge region will tell us about the age of bulge clusters if these exist. In this article we have observed a relative absence of correlation in the bulge that is somewhat less than the prediction of our simple model, but  at best the prediction could say, as in the case of the anticentre, whether the correlation is greater or less than the improved model and enable us to reach further conclusions. \\subsection*"
    },
    "9804/astro-ph9804164_arXiv.txt": {
        "abstract": "We analyze ultraviolet ($\\sim 1500$~\\AA) images of the old open clusters M67, NGC 188, and NGC 6791 obtained with Ultraviolet Imaging Telescope (UIT) during the second flight of the {\\em Astro} observatory in March 1995. Twenty stars are detected on the UIT image of M67, including 11 blue stragglers, seven white dwarf candidates, and the yellow giant -- white dwarf binary S1040. The ultraviolet photometry of the blue stragglers F90 (S975) and F131 (S1082) suggests that these stars have hot subluminous companions. We present a semi-empirical integrated ultraviolet spectrum of M67, and show that the blue stragglers dominate the integrated spectrum at wavelengths shorter than 2600~\\AA. The number of white dwarfs in M67 is roughly consistent with the number expected from white dwarf cooling models. Eight candidate sdB/sdO stars are detected in NGC 6791, and two are detected in NGC 188.    The luminosity range $1.10 < \\log$ L/\\lsun\\ $< 1.27$, derived from the ultraviolet photometry of  the six sdB candidates, is consistent with theoretical models of metal-rich hot horizontal branch (HB) stars.  The fraction of hot HB stars in both NGC 6791 and NGC 188 is about 30\\%, implying that the integrated spectra of both clusters should show a UV turnup at least as strong as that  observed in any elliptical galaxy. ",
        "introduction": "As a stellar population ages, the main-sequence turnoff becomes cooler and fainter, until, after an age of about 2 Gyrs, the ultraviolet  ($\\sim$ 1500~\\AA) flux is no longer dominated by stars near the main-sequence, but instead must be due to a more exotic  population of hot stars, such as blue stragglers, hot subdwarf (sdB, sdO) stars, hot post-asymptotic giant branch (post-AGB) stars, white dwarfs, or interacting binaries.    The blue stragglers and hot subdwarfs are of particular interest since they might have significant roles, respectively, in two current problems in extragalactic astronomy:  the age-dating of old galaxies from their rest-frame near-UV ($\\sim 2600$~\\AA) spectra,  and understanding the origin of the observed UV-upturn in elliptical galaxies.      In the absence of blue stragglers, the rest-frame near-UV spectrum of an old galaxy should be dominated by the turnoff population, and thus fitting the spectrum to population synthesis models should be a reliable method for age-dating the galaxy, although metallicity effects also need to be properly disentangled (\\cite{heap98}, \\cite{bruz97}).       However, as noted by \\markcite{spin97}Spinrad et al. (1997), a significant underestimate of the galaxy age could result if blue stragglers  contribute a large fraction of the galaxy near-UV flux, but are not included in the adopted population synthesis model.  \\markcite{greg90}Greggio \\& Renzini (1990) outlined the various candidates for the origin of the observed UV-upturn in elliptical galaxies, which is characterized by a rising UV flux shortward of 1800~\\AA\\ (\\cite{bica96b}, \\cite{brown97}).   Among these candidates, white dwarfs and hot post-AGB stars have the advantage of being relatively well-understood phases of stellar evolution, but lack sufficient fuel to explain the strongest UV-upturn galaxies.   The more promising candidates are currently the extreme horizontal-branch (EHB) stars and their hot progeny; the AGB-manqu\\'e stars, which miss the AGB phase entirely, and the post-early AGB stars, which leave the AGB before the thermally pulsing phase.     Although the horizontal branch (HB) morphology of the globular clusters becomes generally redder with increasing metallicity, a significant EHB population is theoretically expected at metallicities above solar (and an age larger than about 5 Gyr), if reasonable assumptions are made concerning the metallicity dependence of the mass loss rate and helium abundance  (\\cite{dorm95}, \\cite{yi97a}). Spectroscopically, EHB and AGB-manqu\\'e stars are expected to appear, respectively, as  sdB and sdO stars, and the measured effective temperatures and gravities of the  field hot subdwarf stars is consistent with this scenario (\\cite{saff94}).   However, the field hot subdwarfs apparently have a large binary fraction (\\cite{theis95}, \\cite{green98b}), which instead suggests that binary interactions might be involved in their formation (\\cite{bailyn95}).  Despite their probable contributions to the ultraviolet spectra of galaxies, blue stragglers and hot subdwarfs  are either neglected or poorly constrained in existing population synthesis models of old stellar populations.    The currently  favored explanations for the origin  of blue stragglers involve binary interactions, either through mass-transfer processes,  binary mergers, or direct stellar collisions  (\\cite{bailyn95}, \\cite{leonard96}), but the large parameter space of these binary  processes limits their inclusion in synthesis models to special cases (e.g.\\ \\cite{pols94}).    Some attempts have been made to include EHB stars in synthesis models of metal-rich populations, but these are  hampered by  large uncertainties, for example, in the dependence of mass loss and helium abundance on metallicity (\\cite{bress94}, \\cite{yi97b}). The modeling of galaxy ultraviolet spectra should clearly benefit from empirical studies of the ultraviolet content of resolved, old,  metal-rich stellar populations.  Here we report on ultraviolet ($\\sim$ 1600~\\AA) images of the three old ($> 4$ Gyr) open  clusters M67, NGC 188, and NGC 6791, which were obtained in March 1995  using the {\\em Ultraviolet Imaging Telescope} (UIT, \\cite{stech97}) during the second flight of the {\\em Astro} observatory.      The $40'$ diameter field of  view of UIT is sufficiently large to allow a complete  census of the hot stellar population in the observed clusters.   Thus, the UIT images of these three clusters can be used for an  empirical study of the ultraviolet content of old, metal-rich stellar populations.     As it turns out, M67 contains blue stragglers, but no hot subdwarfs, whereas NGC 188 and NGC 6791 contain hot subdwarfs, but  are too old, distant and reddened for the blue stragglers to be detectable with UIT.    M67 is also sufficiently nearby that all the white dwarfs hotter than $\\sim$ 21,000 K should have been detected on the UIT image.   In the population synthesis models of \\markcite{magris93}Magris \\& Bruzual (1993),  white dwarfs supply about 10\\% of the 1750~\\AA\\ flux in elliptical galaxies with a weak UV-upturn,  and in Section 3.2 we re-examine the white dwarf contribution to galaxy ultraviolet spectra, using more recent stellar atmospheres and evolutionary models.  A brief summary of this work was given by \\markcite{last97}Landsman \\& Stecher (1997). A subsequent paper will discuss the UIT observations of the intermediate age ($\\sim$ 2 Gyr) clusters NGC 752 and NGC 7789, which are sufficiently young to allow the turnoff population to be detected in the ultraviolet. ",
        "conclusions": "There are at least three ways in which Galactic open clusters are an inadequate model for  understanding the stellar populations of old galaxies.    First, the open clusters have experienced dynamical interactions which can alter the stellar population mix; for example, by leading to evaporation of low-mass stars while concentrating the blue stragglers toward the cluster center. Second, the higher mean stellar density of open clusters might result in enhanced stellar interactions and possibly increased blue straggler or hot subdwarf formation. Finally, due to the relatively small population of  open clusters, the rapid evolutionary phases may not be adequately sampled.  For example, hot post-AGB stars are believed to be a significant (although not the dominant)  contributor to the observed UV-upturn in elliptical galaxies, but consistent with their short ($\\sim 10^5$ yr) lifetime, there are no post-AGB stars present in M67, NGC 188, or NGC 6791.    And while hot subdwarfs have now been detected in NGC 188 and NGC 6791, their total number ($<8$) remains uncomfortably low to draw any robust conclusions.  Keeping this limitations in mind, we summarize below the main results derived from the UIT observations of M67, NGC 188, and NGC 6791, and their implications for the study of the integrated ultraviolet spectra of old stellar populations.  \\begin{enumerate}  \\item  The UIT image of M67 is dominated by the eleven detected blue stragglers; in particular, the blue straggler F81 contributes 90\\% of the integrated flux of M67 at 1520~\\AA.    The ultraviolet flux of the two blue stragglers,  F90 and F131, is significantly higher than that predicted on the basis of optical photometry, and probably indicates the presence of hot, subluminous companions.   A semi-empirical calculation of the integrated ultraviolet spectrum of M67 shows that,  even when neglecting the possibly anomalous star F81, the blue stragglers dominate at wavelengths shorter than $\\sim$ 2600~\\AA.   As pointed out by Spinrad et al.\\ (1997), neglect of blue stragglers in population synthesis model fits of the rest-frame near-ultraviolet spectra of a galaxy will result in an underestimate of the galaxy age.  \\item Eight white dwarf candidates are identified in the UIT image of M67, including the core-helium white dwarf companion of the yellow giant S1040. Optical spectroscopy of two of these sources has been obtained by \\markcite{flem97}Fleming et al.\\ (1997): G152 is a hot (\\teff $\\sim$ 68,000 K) DA white dwarf, and MMJ 5973 is cooler (\\teff\\ $\\sim$ 18,000 K) DB white dwarf. The number of white dwarf candidates is in reasonable agreement with that expected from theoretical white dwarf cooling models and a cluster age of 4 Gyr.    The integrated ultraviolet flux at 1500~\\AA\\ along a white dwarf cooling model is $\\log$ E$_{1500} = -4.23$ L$_{V}^{\\odot}$ Gyr \\AA$^{-1}$, and the contribution of white dwarfs to the integrated spectra of old galaxies is roughly 10\\% of that expected from hot post-AGB stars.  \\item Eight probable cluster members are detected on the UIT image of NGC 6791, including the five hot subdwarfs studied spectroscopically by Liebert et al.\\ (1994), and the two additional sdB/O candidates, B9 and B10, reported by \\markcite{kal95}Kaluzny \\& Rucinski (1995).    Three probable cluster members are detected on the UIT image of NGC 188, including the sdB spectroscopic binary II-91.   The star D702 in NGC 188 is probably a composite including a hot subdwarf, since it has a hot ultraviolet color (\\mbi\\ -- \\mbv\\ = --0.11), but a relatively cool optical color (\\bv\\ = 0.26).  The derived  luminosity range, $1.10 < \\log$ L/\\lsun\\ $< 1.27$, of the five sdB stars in NGC 6791, and II-91 in NGC 188, is consistent with that expected for metal-rich, hot HB stars.    The fraction of hot HB stars in both clusters is about 30\\%, implying that the integrated spectra of both clusters should show a pronounced UV-upturn, as strong as that observed in any elliptical galaxy.   \\end{enumerate}"
    },
    "9804/astro-ph9804328_arXiv.txt": {
        "abstract": "Galaxy distance indicators are subject to different uncertainties and biases which may reflect in the peculiar velocity field. In this work, we present different statistical analysis applied to peculiar velocities of samples of galaxies and mock catalogs taken from numerical simulations considering these observational effects. Our statistical studies comprise the mean relative velocity and velocity correlation for pairs of galaxies as a function of separation, and a bulk flow analysis determined on spheres of $10 h^{-1}$ $Mpc$ radius. In order to test the statistical analysis we use COBE normalized CDM numerical simulations with different density parameters and cosmological constant and take into account the Tully-Fisher (TF) scatter and a possible TF zero-point offset, as well as variations in the results of the simulations arising from different observer positions. We compare the results of the mock catalogs with samples of spiral galaxies taken from the Mark III catalog. The results of our analysis indicate the importance of errors in deriving the density parameter $\\Omega$ from statistics applied to the observational data and it is found that large values $\\Omega \\geq 0.8$ may be derived from the analysis if errors are not taken into account. However, the observed value of the TF scatter ($\\simeq 0.35$ $mag$) make CDM models with $\\Omega > 0.6$, inconsistent with the statistical tests involving relative velocities. A suitable agreement is found for models with $\\Omega \\leq 0.6$, although requiring a TF zero-point offset of the order of a tenth of a magnitude to provide consistency with the observed flow coherence. ",
        "introduction": "Measurements of peculiar velocities provide direct probes of the mass distribution in the Universe putting constraints on models of large-scale structure formation (e.g. \\cite{peebles}, \\cite{vittorio}). Advances on new distance indicators (\\cite{djor}; \\cite{dresslera}) have allowed estimates of peculiar flows in the local Universe up to $\\sim 50-100 h^{-1}$Mpc (see \\cite{giova}, or \\cite{strauss} for a review). The results from using elliptical galaxies (\\cite{dresslerb}) suggest a large coherence length and amplitude of the local peculiar velocity field. The results from analysis of 1,355 spirals using the Tully-Fisher relation to determine distances (\\cite{mathew}) would lead to an even larger coherence length and a similar amplitude. Other samples are also in agreement with the Great Attractor findings (\\cite{willick}). In a more controversial finding \\cite{postman} find from analysis of clusters of galaxies that the peculiar flows can be coherent over a much larger scale than the Great Attractor findings. All the results indicate significant deviations from the Hubble flow with a large coherence length.  Since peculiar velocities probe the mass distribution and also depend on the density parameter $\\Omega$ one can determine the latter by comparing the mass distribution implied by the velocity field with the observed distribution of galaxies. The density parameter is thus determined to within the uncertainty of the bias factor $b$ by determining the factor $\\Omega^{0.6}/b$. Other analysis of the peculiar velocity information may be applied in order to obtain the parameter $\\Omega$. \\cite{berta} developed the POTENT method whereby the mass distribution is reconstructed by using the analog of the Bernoulli equation for irrotational flows. They used the method to analyze to great detail the peculiar velocity field out to about 60$h^{-1}$Mpc (\\cite{bertb}). \\cite{dekel} compared the previously determined velocity field with the observed distribution of galaxies and concluded that the best fit is with $\\Omega^{0.6}/b \\simeq 1$. A similar analysis can be applied to the components of the velocity tensor $U_{ij}$ (\\cite{gorskia}; \\cite{groth}). Recently, \\cite{kashlinsky} has analyzed this tensor for the MarkIII peculiar velocity data to reconstruct the large-scale power spectrum obtaining consistency within the CDM model with $\\sigma_{8} \\Omega^{0.6} \\simeq 0.8$. In a similar study \\cite{zaroubi} obtain a different result $\\sigma_{8} \\Omega^{0.6} \\simeq 0.35$ indicating the need of further analysis.   In this paper we study the peculiar velocity field through a statistical analysis of COBE normalized CDM numerical simulations and observational data taken from the Mark III catalog. Our comparison of models and observations take into account observational uncertainties and possible biases, and variations of the results according to different observers in fully non-linear numerical simulations.    \\section {Data}  We use the Mark III catalog (\\cite{markIII1};\\cite{markIII2};\\cite{markIII3}) as a suitable data set to analyze the peculiar velocity flow. This Catalog lists Tully-Fisher and $Dn-\\sigma$ distances and radial velocities for spiral, irregular, and elliptical galaxies. In our analysis we have used only spiral galaxies given their large number and smooth spatial distribution (see Table 1).  \\begin{deluxetable}{cccc} \\tablewidth{35pc} \\tablecaption{Observations: The Mark III spirals} \\tablehead{ \\colhead{Subsample}& \\colhead{$N^o$ of Gx.}& \\colhead{TF relation}& \\colhead{$\\sigma_{TF}$} } \\startdata \\phm{II} Aaronson et al. Field & 359 &  $M_H=-5.95+10.29 \\eta$ & $0.47$ \\nl \\phm{II} Mathewson et al. (1992) & 1355 & $M_I=-5.79+6.8 \\eta$ & $0.43$ \\nl \\phm{II} Willick, Perseus Pisces (1991) & 383 & $M_r=-4.28+7.12 \\eta$ & 0.38 \\nl \\phm{II} Willick, Cluster Galaxy (1991) & 156 & $M_r=-4.18+7.73 \\eta$ & $0.38$ \\nl \\phm{II} Courteau-Faber (1993) & 326 & $M_r=-4.22+7.73 \\eta$ & $0.38$\\nl \\phm{II} Han-Mould et al., Cl. Gx. (1992) & 433 & $M_I=-5.48+7.87 \\eta$ & 0.4 \\nl \\enddata \\end{deluxetable}   The velocity parameter $\\eta = Log \\Delta V - 2.5$ is determined either from HI profiles or from optical $H_{\\alpha}$ rotation curves. The TF relations and their corresponding scatters for the different samples of spiral galaxies are given by \\cite{markIII3} and are shown in Table 1, where the absolute magnitud $M$ satisfies $M=m-5\\log cz$. The galaxy apparent magnitudes $m$ of the Tully-Fisher distances are corrected for Galactic extinction, inclination and redshift (see Willick et al. 1997 for details).  The selection bias in the calibration of the forward TF relation can be corrected once the selection function is known. But then, the TF inferred distances and the mean peculiar velocities suffer from Malmquist bias. Suitable  procedures to consider these biases, induced both by inhomogeneities and selection function, have been discussed (see for instance /cite{freud95} and references therein) where the spatial distribution, selection effects and observational uncertainties are realistically modeled through Monte-Carlo simulations. We have used in our analysis forward TF distances, fully corrected for Inhomogeneous Malmquist Bias (\\cite{markIII1}, \\cite{markIII2}, \\cite{markIII3}). Nevertheless we have found that the results of the statistics studied in this work do not change significantly if inverse TF distances are used as it will be discussed below.  Radial velocities used to infer the peculiar velocity of the galaxies are referred to the Cosmic Microwave Background frame. It should be remarked that galaxy distance estimates are subject to errors due to the scatter in the TF relation (\\cite{mo}; \\cite{willickth}; \\cite{mathew}) and uncertainties of the TF zero-point (\\cite{shanks}; \\cite{willickth}). Also, the possible presence of a small fraction of spurious velocities in the data induced by either galaxy peculiarities or observational errors in distance estimates (\\cite{jacoby}) should be taken into account. ",
        "conclusions": "In this work we have attempted to asses the effects of galaxy distance uncertainties and biases on statistical tests of the peculiar velocity field using samples of spiral galaxies from the Mark III compilation and mock catalogs taken from numerical simulations. Our studies comprise relative velocity tests and pair velocity correlations as a function of separation, as well as a bulk flow analysis determined on spheres of $10 h^{-1}$ $Mpc$ radius. We construct mock catalogs using numerical simulations corresponding to COBE normalized CDM models with different values of density parameter and cosmological constant. The models take into account the Tully-Fisher (TF) scatter and possible TF zero-point shift, as well as variations in the results of the simulations arising from different observer positions. The analysis of the departures from gaussianity of the observational velocity distributions as measured by the kurtosis tests $K$ and $K1$ show that only a small fraction ($<2 \\%)$ of possible spuriously high velocities (and therefore biased distance estimates) might be present in the observational data. We find that the observed scatter of the TF relation plays a very important role when deriving constraints to the cosmological parameters. If errors in galaxy distance estimates are neglected, the observed magnitude of the $\\Delta V$ results show consistency with high density CDM models ($\\Omega>0.8$). When the observed TF scatter $\\Delta M \\simeq 0.35$ is taken into account a significant disagreement with observations is found for the $\\Delta V$ and $D1$ statistics in models with $\\Omega > 0.6$.  A possible offset in the TF zero-point used to determine the distances of the galaxies in the observational data artificially enhance velocity correlations ($\\Pi$ and $M$ statistics). Due to this fact, if zero-point offsets of the order of a tenth of a magnitude are considered, the  CDM models explored provide a more satisfactory fit to observations in these tests.  The results presented in this work are not sensitive to sample variations. We have applied the same statistical tests to the subsample corresponding to Mathewson data and we find similar results than those of the total sample. It should be noted that the variations on the statistics arising from different observers are significantly enhanced when errors on the peculiar velocities are included.  This fact make more difficult the distinction between the models. Error bars in figures 1 to 5 show the variations of the results arising from different observer positions in the models and serve as a test for the dependence of the observational results on our particular position in space. In our analysis of biased models we have not found relevant differences in the results of the statistical analysis of samples with different density thresholds. The selection of low (high) density particles in the simulations does not produce very relevant effects in the statistics although lower (higher) values of $\\Pi$ and $\\Delta V$ are observed due to the oversampling (undersampling) of high density particles where random motions dominate. Finally when comparing models and observations, we find that neither the K, K1, M, nor $\\Pi$ statistical tests are sensitive to the density parameter of the models.     We find crucial to consider properly the intrinsic scatter of the Tully-Fisher relation in studies of the peculiar velocity field. Although low values of the density parameter $\\Omega < 0.6 $, are favored by our statistical analysis  $D1$ and $\\Delta V$ when this scatter is taken into account $\\pm 0.15$ $mag$ zero-point offsets in the TF relation are required in order to provide the observed values of $\\Pi$ and $M$ statistics in the models explored. A positive TF zero-point offset of this magnitude would imply lower values of the Hubble constant which may be worth to consider in a controversial  topic (see for instance \\cite{shanks} and \\cite{th} and references therein)."
    },
    "9804/astro-ph9804091_arXiv.txt": {
        "abstract": "The conversion of neutron matter into strange matter in a neutron star occurs through the non-leptonic weak-interaction process. We study the energy loss of the neutron star by the emission of axions in that process. Owing to that process, the neutron star will liberate the energy which can in no way be negligible as an axion burst. \\\\ \\\\ PACS numbers : 14.80.Mz, 95.30.Cq, 97.60.Jd \\\\ Keywords : axions, astrophysics, neutron stars, strange star \\\\ \\\\ ",
        "introduction": "The possibility of the existence of stable quark matter in the early universe or inside a neutron star  or in relativistic heavy-ion collision experiments has been studied in many works \\cite{sqm}\\@. In this work, we consider the conversion of a neutron star into a strange star and its energy loss during the conversion process. In order to form strange matter in the interior of a neutron star, neutron matter should be converted into strange matter. Several conversion mechanisms have been discussed by Alcock et al.\\ \\cite{alcock}\\@. When a neutron and a stable strangelet (strange quark matter droplet) meet, the neutron is readily absorbed, while a proton can coexist with a strangelet due to the Coulomb barrier between them \\cite{alcock}\\@. Therefore, a stable strangelet as a seed for the process will trigger a conversion in such a way that it grows by absorbing neutrons and, eventually, convert most of the neutron star into a strange star \\cite{olinto,daietal}\\@. Then the conversion would liberate about $10^{52}$ ergs in binding energy \\cite{olinto}\\@.  A variety of mechanisms have been suggested for seeding the interior of a neutron star with stable strangelets so far \\cite{alcock,lugones}\\@. They can be divided into two main categories \\cite{olinto}; (1) the primary mechanisms in which the seed is formed inside the neutron star, (2) the secondary mechanisms in which the seed comes from the interstellar medium. In a viewpoint of the primary mechanisms \\cite{alcock}, the central high densities and pressures in a neutron star may be sufficient for a phase transition from neutron matter to two-flavor quark matter to occur. Subsequently two-flavor quark matter can then easily decay into the lower energy strange matter through weak interactions.  If we assume that strangelet has formed inside a neutron star, subsequent conversion process of the rest of the star is described as follows \\cite{olinto,olesen}\\@. The volume over which strange matter equilibrates was shown to be much smaller than that of the total strange matter region \\cite{lugones,horvath}, so that the problem can be treated hydrodynamically in one-dimensional geometry. As the conversion front sweeps into neutron matter, the small region behind the conversion front has an excess of down quarks relative to strange quarks due to the flux of neutrons ($udd$) at the conversion front. The excess down quarks will convert into strange quarks via the non-leptonic weak process \\cite{madsen}, $d \\; + \\; u \\; \\rightarrow \\; s \\; + \\; u$, as long as $\\mu_d > \\mu_s (n_d > n_s)$\\@. By this process a $d$ quark can change itself to an $s$ quark until the Fermi energies of all the flavors become equal and the energy per baryon drops in comparison to the ordinary two-component nuclear matter. Other leptonic decay processes are suppressed considerably \\cite{madsen}\\@. The conversion will liberate $\\sim 10^{52}$ ergs of energy (assuming that $\\sim$ 10 $MeV$ is liberated per neutron converted)\\@. This energy will be radiated as neutrinos, photons, $e^{+}e^{-}$ pairs, and axions, etc. In this stage neutrinos and axions can escape the star. Many authors \\cite{daietal,aggs} have considered the escape of neutrinos. In this work we calculate the energy loss due to the emission of axions in the process of such a conversion of a neutron star into a strange star.  At the chemical equilibrium between the quarks and the electrons, non- or semi-leptonic weak interaction is not important in neutrino and axion emission in quark matter because the weak interactions coupling $s$ and $u$ quarks are Cabibbo suppressed relative to the interactions coupling $d$ and $u$ quarks. For the nonequilibrium processes such as the conversion from two-flavor to three-flavor quark matter, however, as the strange quark semi-leptonic processes in neutrino emissions and non-leptonic processes in axion emissions are of significance.   The axion is a pseudo-Goldstone boson which was introduced by Peccei and Quinn (PQ) to solve the strong CP-problem in a natural way \\cite{PQ}\\@. However, theoretical and experimental investigations give little guidance on the PQ symmetry-breaking scale, $f_a$, and therefore on the mass of the axion \\cite{WW}\\@. The axion decay constant $f_a$ is related to the axion mass \\cite{axion} \\[ m_a \\simeq \\left( \\frac{0.62 \\times 10^{7} GeV}{f_a} \\right) \\; eV . \\] There are two generic types of invisible axions; the Dine-Fischler-Srednicki-Zhitnitskii (DFSZ) axion which couples to both quarks and leptons at tree level \\cite{DFSZ} and the hadronic, or Kim-Shifman-Vainshtein-Zakharov (KSVZ), axion which has no tree-level coupling to leptons but does couple to them at loop levels \\cite{KSVZ}\\@. In invisible axion models, the axion mass $m_a$ is in principle arbitrary, however astrophysical and cosmological considerations \\cite{axion} can provide an upper and lower bounds for $m_a$\\@. The astrophysical bounds on $m_a$ are due to the fact that axion emission is an additional energy loss mechanism for stars.  If such axions could be copiously produced during the conversion of a neutron star into a strange star, it might drastically alter the energy budget of stars. The axions might thus deplete the stellar energy and change the usual course of stellar evolution. The emission of axions would have hastened the cooling process. The quarks are expected to be highly relativistic and degenerate in such matter with their Fermi energies in the range of 300 to 500 $MeV$\\@. A calculation of the axion energy flux from strange quark matter has been made by Anand et al.\\ \\cite{anand}\\@. They have shown that the axion emission rate is several orders of magnitude smaller compare to the neutrino emission rate.  We investigate here the energy loss of the neutron star during the conversion of non-strange quark matter into strange quark matter in the interior of the star. We assume that the axions escape freely out of the star as soon as they are produced, and that the temperature in the interior is constant during the conversion for the reasons that strong and electromagnetic interaction timescales are much smaller than those of weak interactions. Unless otherwise noted all equations assume $\\hbar = c = k_B = 1$\\@. ",
        "conclusions": "Our main result for the total energy flux rate for the emission of axions from conversion of non-strange quark matter into strange quark matter is given by Eq. (12). Here we can carefully quote our bound on the axion coupling $f_a$, using the simple inequality \\cite{ishizuka}, \\be {\\cal E}_a \\cdot V \\cdot \\delta t < E_m \\ee The allowed energy loss, $E_m / \\delta t$ is taken to be $10^{52}$ erg/sec for definiteness, and the volume of axion emitting region $V = \\frac{4}{3} \\pi R^3$ with $R = 10 \\sim 20 km$\\@. As for the conversion timescale, if $T=10 MeV$, it varies roughly between one and ten minutes \\cite{olinto}\\@. So, we take the timescale typically $\\delta t = 100 sec$\\@. Then, from Eq. (14), we can obtain the bound on $f_a$, \\be \\sim 10^5 \\; GeV < f_a \\,. \\ee We note that the bound based on Eq. (14) is legitimate only if the emitted axions never interact on their way out of the star. If the axion coupling is large enough, axions once produced may interact many times, namely be absorbed and reemitted more than once in the hot core of the star, and may be trapped thermally.  As applications, Mikheev and Vassilevskaya \\cite{mik} recently investigated the radiative decay of the axion $a \\rightarrow \\gamma\\gamma$ in an external electromagnetic field in the DFSZ model. They concluded that the decay probability is strongly catalyzed by the external field, namely, the field removes the main suppression caused by the smallness of the axion mass. Therefore, the radiative decay of the axion in strong magnetic fields of the neutron star could give interesting astrophysical phenomena. Furthermore, we can consider that the sudden conversion from a neutron star to a strange star may account for the gamma ray bursts at comological distances \\cite{olinto}\\@. Most of the energy is probably released in the form of neutrinos. If a part of the total energy goes into $\\gamma$ rays decayed from the axions, it will be large enough to account for the gamma ray bursts at cosmological distances and to explain their isotropic distribution. The outcome of such a conversion event will be the emission of as much as $\\sim 10^{58} MeV$ of energy in $\\sim sec$ to $\\sim min$ as radiation with a typical temperature of tens of $MeV$\\@. This conversion can be observed as a gamma ray burst \\cite{olinto}\\@.  In summary, as a possible mechanism for production of axions, we considered the conversion of non-strange quark matter into strange quark matter. We also estimated the energy loss of neutron stars through the emission of axions in addition to the cooling provided by the neutrino emission. During the conversion period, the important process is the non-leptonic weak-interaction. We assume that the axions are not trapped and escape freely out of the star as soon as they are produced, and that the temperature in the interior is constant during the conversion. We find that the energy carried away by axions during the conversion is not negligible. It is also found that the axion emission rate is three orders of magnitude larger compared to the chemical equilibrium case. In addition, we discussed the bound on the axion coupling.  \\vspace*{3cm}"
    },
    "9804/astro-ph9804058_arXiv.txt": {
        "abstract": "We have analyzed deep $B$ and $V$ photometry of the Carina dwarf spheroidal reaching below the old main-sequence turnoff to $V \\sim 25$.  Using simulated color-magnitude diagrams to model a range of star formation scenarios, we have extracted a detailed, global star formation history.  Carina experienced three significant episodes of star formation at $\\sim 15$ Gyr, 7 Gyr, and 3 Gyr.  Contrary to the generic picture of galaxy evolution, however, the bulk of star formation, at least 50\\%, occured during the episode 7 Gyr ago, which may have lasted as long as 2 Gyr.  For unknown reasons, Carina formed only 10-20\\% of its stars at an ancient epoch and then remained quiescent for more than 4 Gyr.  The remainder ($\\sim 30\\%$) formed relatively recently, only 3 Gyr ago. Interest in the local population of dwarf galaxies has increased lately due to their potential importance in the understanding of faint galaxy counts.  We surmise that objects like Carina, which exhibits the most extreme episodic behavior of any of the dwarf spheroidal companions to the Galaxy, are capable of contributing to the observed excess of blue galaxies at $B \\sim 24$ only if the star formation occurred instantaneously. ",
        "introduction": "The Carina dwarf galaxy, one of the nine known dwarf spheroidal companions to the Galaxy, has a surprisingly complex star formation history.  Detection of carbon stars provided the first suggestion of a significant intermediate-age population (Cannon, Niss, \\& Norgard-Nielsen 1981), followed by main-sequence photometry which revealed a young turn-off due to a population perhaps only 6--9 Gyr old (Mould \\& Aaronson 1983; MA hereafter).  MA fit a simulated luminosity function comprised of 7 Gyr-old stars to their data, and estimated the contribution of an old population to be relatively small in comparison.  Saha, Monet, and Seitzer (1986) observed a large number of RR Lyraes in the central region of the galaxy and established a lower limit of 2--3\\% for the fraction of old stars, depending on the yield of RR Lyraes. Mighell (1990) estimated the relative sizes of the populations from the double-peaked color distribution near the MSTO region to be 85\\% for the intermediate-age burst and 15\\% for the old episode.  Using simulated luminosity functions, Mighell \\& Butcher (1992) later fit intermediate-age burst models to this deep, main-sequence photometry and estimated an upper limit to the old population of 40\\%.  More recently, Smecker-Hane, et al. (1994; hereafter SHSHL) resolved two separate horizontal branches (HBs).  Although the separation is not independent of metallicity effects, it is very suggestive of the multi-episode nature of Carina's history.  We are left with an estimate between 2\\% and 40\\% for the fraction of old stars, and no clear result on the duration of the star formation episodes.  The chemical evolution of the galaxy is likewise still in question.  Spectroscopy of 15 giants in Carina (Da Costa 1994) resulted in an average [Fe/H] of --1.9, with one giant significantly more metal poor ([Fe/H] = --2.2).  The excellent areal coverage of the SHSHL data reveal a very thin giant branch.  Given the huge apparent age spread, a metallicity spread may be necessary to compensate and produce the observed narrow RGB.  Improved methods of analysis are needed.  Detailed analyses of stellar luminosity functions have some advantages over more traditional isochrone fitting.  These methods, however, were designed for coeval systems.  Because dwarfs exhibit a range of ages and metallicities, these approaches are problematic at best, and misleading at worst.  A much better way to unravel complex star formation histories like those exhibited by dwarf galaxies such as Carina and Leo~I is to use all of the information embodied by color-magnitude diagrams.  In these cases, where the galaxies appear to have experienced bursts at intermediate epochs, a conventional luminosity function would not distinguish between the old subgiant branch and the young MS stars, for example. Color information proves essential to resolve multiple components in these systems.  Simulated color-magnitude diagrams have been used successfully in studies of bright stars in dwarf irregulars (Tosi et al. 1991; Tolstoy 1996; Gallart et al. 1996a, 1996b, Aparicio et al. 1997a, 1997b), of LMC field stars (Bertelli et al. 1992; Vallenari et al. 1996), and of LMC clusters (Vallenari et al. 1994).  We have deep CCD photometry (limiting magnitude $V = 24.5$) with reasonably small photometric errors ($\\sim 0.02$ at $V=23$) of three fields in the Carina dwarf galaxy, reaching well below the old main-sequence turnoff.  Because the photometry reaches below the old MSTO, these data are well-suited to a detailed extraction of the star formation history by comparing model color-magnitude diagrams to that of Carina.  Section 2 is a discussion of the observations and reductions of the Carina data.  Section 3 outlines the analysis applied to these data, including the development of the pseudo-LF and the generation of similulated color-magnitude diagrams.  The results, presented in section 4, are summarized and briefly discussed in the context of galaxy evolution in section 5. ",
        "conclusions": "Carina experienced an episode of star formation 7 Gyr ago which lasted no more than 2 Gyr and which was responsible for at least 50\\% of its stars.  The old population ($12-15$ Gyr) may amount to $10\\%-20\\%$ of Carina. The bulk of the remainder (($20\\%-30\\%$) is relatively young, between 2.5 and 3.5 Gyr old.  While the details of one galaxy's SF history may seem inconsequential, taken in context and as a member of the Local Group dwarf population, the details are relevant to deeper, unresolved issues in cosmology and galaxy evolution.  In the case of Carina, all studies confirm that at least one relatively long pause in star formation lasting $\\sim 4$ Gyr occurred. The mechanism by which this kind of low mass system could experience such a pause and then another strong burst of star formation is not understood.  Current ISM simulations can produce such a large gap in an isolated dwarf only by ejecting the gas out to $\\sim 20$ kpc (Babul 1996).  A dwarf such as Carina, residing in the potential well of the Galaxy, should lose that gas, preventing any further episodes of star formation.  Recent Hubble Space Telescope observations of Carina have been studied by Mighell (1997). He claims to detect a significant number of stars in the region of the gap in star formation lasting from roughly 8 to 12 Gyr ago.  This is interpreted as indicative of significant star formation which began in the central region of Carina, and propagated to the outer regions.  However, the number of stars in the same section of our CMD, derived from fields several arcminutes from the center of Carina and overlapping the Mighell (1990) field, is consistent with the number in the HST CMD.  Whether or not the number of stars in this area implies significant on-going star formation, there does not seem to be evidence of a difference between the central region and regions further out, such as that seen in the recently-discovered Antlia dwarf galaxy (Aparicio et al 1997c).  Assigning ages to individual faint stars based on the isochrones which constrain them in color and magnitude must be handled with care. Firstly, at faint magnitudes, stars have a larger photometric error and thus a larger implied age range.  Age determinations for specific stars are therefore subject to greater statistical uncertainty. Secondly, the lifetime of stars crossing the Hertzprung gap increases as stars become less massive.  This must be accounted for when predicting the relative contributions of different ages. In light of these considerations, the results of the WFPC2 study may not differ from earlier ground based results.  Evidence of intermediate-age stars in the halo leads to the question of whether dwarf spheroidals shredded by the Galaxy could be a significant source of the halo population (Preston et al. 1994; Mateo 1996).  The Sagittarius dSph is currently being ripped apart by the Milky Way, proving that this type of interaction between a large galaxy and its tiny neighbors does occur (Ibata, Gilmore, \\& Irwin 1994; Mateo et al. 1995).  In addition, the dSph, especially Carina and Leo I, contain significant numbers of intermediate-age stars. Mateo (1996) shows that the fraction of relatively young and intermediate-age stars in the entire population of dwarf spheroidals is not inconsistent with the fraction in the Galactic halo.  The detailed interpretation of deep galaxy counts and redshift surveys depends on the SFH of the dwarfs which comprise the excess of faint galaxies.  For example, `flashing' dwarfs - i.e. bursting quickly and intensely - contribute to the deep counts to a dramatically different extent than do dwarfs that experience more drawn-out bursts that slowly turn on and off (Campos 1997).  The prevalence of bursting behavior in the SFH of Local Group dwarfs may statistically constrain the degree to which galaxies similar to these local dwarfs could be responsible for the faint excess.  We can place an upper limit on the star formation rate (SFR) of Carina 7 Gyr ago by assuming that the episode of star formation took place instantaneously.  By choosing a reasonable model, we can easily calculate the total number of stars formed during this episode in the volume of the galaxy covered by our three fields, taking into account the incompleteness.  We can use this number to normalize the IMF, and then integrate to get the total mass formed during this episode.  The SFR is then given by: \\begin{displaymath} SFR = \\frac{M_{total}}{\\Delta t}*\\frac{1}{\\gamma   }  M_{\\odot  }  yr^{-1} \\end{displaymath} where $\\Delta t$ is the duration of the episode in years and $\\gamma$ is the fraction of the galaxy's total surface brightness in the three fields.  If the episode lasted 10 Myr, the implied SFR for Carina at that time would be roughly 0.8 $M_{\\odot  } yr^{-1}$.  If the episode lasted 2 Gyr, the SFR is down by more than a factor of 100 to 0.004 $M_{\\odot  } yr^{-1}$.  Either SFR is comparable to the instantaneous SFR in even such luminous objects as blue compact dwarfs (Fanelli et al. 1988).  Could Carina-type galaxies contribute to the excess of faint galaxies at around $B \\sim 24$?  These galaxies have been shown to have redshifts between 0.3 and 0.7 (Glazebrook et al. 1995).  Assuming $H_{\\circ}$ = 50 km/s/Mpc, and $\\Omega$ = 1, 7 Gyr corresponds to $z \\sim 0.5$, which places the episode of star formation in the appropriate epoch.  If the burst were essentially instantaneous, the implied total luminosity of Carina at 7 Gyr would be $\\sim 2 \\times 10^8 L_\\odot$.  Carina would be visible to $z \\sim 0.3$ in a sample limited to $V \\sim 25$, putting it at the near edge of objects that could contribute to these counts.  Galaxies 3-10 times more luminous would fall in the range $z \\sim 0.5$ to 1.0.  Although the local dSph cover a range up to 50 times the luminosity of Carina, Carina is the most extreme example of this type of episodic behavior. Further, if the episodes of star formation are extended in time, then these dwarfs would only be visible locally.  Thus, despite concerted efforts at unravelling the SFH of these galaxies, their role in faint galaxy counts problem remains unresolved."
    },
    "9804/astro-ph9804172.txt": {
        "abstract": "We have analyzed the projected galaxy distributions in a subset of the ENACS cluster sample, viz. in those 77 clusters that have z $<$ 0.1 and R$_{\\rm ACO} \\ge$ 1 and for which ENACS and COSMOS data are available.  For 20~\\% of these, the distribution of galaxies in the COSMOS catalogue does not allow a reliable centre position to be determined.  For the other 62 clusters, we first determined the centre and elongation of the galaxy distribution. Subsequently, we made Maximum-Likelihood fits to the distribution of COSMOS galaxies for 4 theoretical profiles, two with `cores' (generalized King- and Hubble-profiles) and two with `cusps' (generalized Navarro, Frenk and White, or NFW, and de~Vaucouleurs profiles).  We obtain average core radii (or characteristic radii for the profiles without core) of 128, 189, 292 and 1582 kpc for fits with King, Hubble, NFW and de~Vaucouleurs profiles respectively, with dispersions around these average values of 88, 116, 191 and 771 kpc. The surface density of background galaxies is about 4 10$^{-5}$ gals arcsec$^{-2}$ (with a spread of about 2 10$^{-5}$), and there is very good agreement between the values found for the 4 profiles. There is also very good agreement on the outer logarithmic slope of the projected galaxy distribution, which is that for the non-generalized King- and Hubble-profile (i.e.  $\\beta_{King}$ = $\\beta_{Hubble}$ = 1, with the corresponding values for the two other model-profiles).  We use the Likelihood ratio to investigate whether the observations are significantly better described by profiles with cusps or by profiles with cores. Taking the King and NFW profiles as `model' of either class, we find that about 75 \\% of the clusters are better fit by the King profile than by the NFW profile.  However, for the individual clusters the preference for the King profile is rarely significant at a confidence level of more than 90 \\%.  When we limit ourselves to the central regions it appears that the signifance increases drastically, with 65 \\% of the clusters showing a strong preference for a King over an NFW profile. At the same time, about 10 \\% of the clusters are clearly better fitted by an NFW profile than by a King profile in their centres.  We constructed composite clusters from the COSMOS and ENACS data, taking special care to avoid the creation of artificial cusps (due to ellipticity), and the destruction of real cusps (due to non-perfect centering). When adding the galaxy distributions to produce a composite cluster, we either applied no scaling of the projected distances, scaling with the core radii of the individual clusters or scaling with r$_{200}$, which is designed to take differences in mass into account. In all three cases we find that the King profile is clearly preferred (at more than 95 \\% confidence) over the NFW profile (over the entire aperture of 5 core-radii). However, this `preference' is not shared by the brightest (M$_{b_j}$ $\\la$ -18.4) galaxies. We conclude that the brighter galaxies are represented almost equally well by King and NFW profiles, but that the distribution of the fainter galaxies clearly shows a core rather than a cusp.  Finally, we compared the outer slope of the galaxy distributions in our clusters with results for model calculations for various choices of fluctuation spectrum and cosmological parameters. We conclude that the observed profile slope indicates a low value for $\\Omega_0$. This is consistent with the direct estimate of $\\Omega_0$ based on the $\\frac ML$-ratios of the individual clusters. ",
        "introduction": "Until fairly recently, the projected galaxy density in rich galaxy clusters was generally described by King or Hubble profiles. In these profiles, the logarithmic slope of the mass distribution is essentially zero near the cluster centre. The core radius which is the characteristic scale of the distribution, was sometimes also regarded as the distance which more or less separates dynamically distinct regions in a cluster. From the kinematics of the galaxy population it appears that in clusters the relaxation time is significantly shorter than the Hubble time {\\em only} in the very central region within at most a few core radii (see e.g. den Hartog and Katgert 1996).  The concept of cores in clusters has been seriously challenged, on observational grounds (e.g. Beers $\\&$ Tonry 1986) and as a result of numerical simulations. Navarro, Frenk and White (1995, 1996) found e.g. that the equilibrium density profiles of dark matter halos in universes with dominant hierarchical clustering all have the same shape, which is essentially independent of the mass of the halo, the spectrum of initial density fluctuations, or the values of the cosmological parameters. This `universal' density profile (NFW profile hereafter) does not have a core, but has a logarithmic slope of --1 near the centre which, at large radii, steepens to --3, and thus closely resembles the Hernquist (1990) profile except for the steeper slope of the latter at large radii of --4.  Navarro, Frenk and White (1997) argue that the apparent variations in profile shape, as reported before, can be understood as being due to differences in the characteristic density (or mass) of the halo, which sets the linear scale at which the transition of the flat central slope to the steep outer slope occurs. They also argued that the existence of giant arcs in clusters requires that the mass distributions in clusters does not exhibit a flat core in the centre. In other words: if clusters have cores, the lensing results require that the core radii are very small, at least quite a bit smaller than the values usually quoted.  It is not clear that galaxy clusters should have cores; after all, the dynamical structure of galaxy clusters is quite different from that of globular clusters, for which Michie \\& Bodenheimer (1963) and King first proposed density profiles with cores, in particular the King profile (see e.g. King 1962). On the other hand, the X-ray data for clusters are quite consistent with the existence of a core in the density distribution. More specifically, it was argued recently by Hughes (1997) that the NFW profile would induce a temperature gradient. The existence of such a gradient in the Coma cluster can be excluded at the 99$\\%$ confidence level. Similarly, the galaxy surface density in clusters is generally found to be consistent with a King profile. For galaxy clusters, little use has been made of the de Vaucouleurs profile to describe the galaxy density, even though the latter was found to arise quite naturally in N-body simulations of the collapse of isolated galaxy systems (e.g. van Albada 1982).  In view of the claimed universality of the NFW profile found in the simulations, it seems useful to have a closer look at the projected distribution of the galaxies in clusters. After all, the NFW profile refers to the total gravitating mass, and it is not obvious that the galaxy distribution should have exactly the same shape as the distribution of total mass; although in numerical experiments no strong biasing between dark and luminous matter in clusters was seen (e.g. van Kampen 1995). In this respect, it is noteworthy that Carlberg et al. (1997) find that the combined galaxy density profile of 16 high-luminosity X-ray clusters at a redshift of $\\approx$0.3 closely follows the NFW profile. More precisely, the logarithmic slope in the central region is consistent with the value of --1 of the NFW and Hernquist profiles, while the outer slope is consistent with both --3 (the NFW value) and --4 (the Hernquist value).  The outer slope of the density profile was found by several authors (e.g. Crone et al. 1994, Jing et al. 1995, and Walter and Klypin 1996) to reflect the details of the formation scenario, and in particular the value of the density parameter of the universe. In addition, this slope is unlikely to be constant in time but is expected to get steeper with decreasing redshift. For that reason, it is important to study both the characteristics of the density profiles of rich clusters and their dependence on redshift.  In this paper, we investigate the projected galaxy distributions for a sample of 62 rich and nearby (z $\\la$ 0.1) clusters. These clusters are taken from the volume-limited ENACS (ESO Nearby Abell Cluster Survey) sample of R$_{\\rm ACO} \\ge 1$ clusters (see e.g. Katgert et al. 1996 (paper I), Mazure et al.  1996 (paper II), Biviano et al. 1997 (paper III), Adami et al. 1998 (paper IV), Katgert et al. 1998 (paper V) and de Theije $\\&$ Katgert 1998 (paper VI)). In $\\S$ 2, we first describe the sample of clusters that we used, and the data on which we based our analysis. In $\\S$ 3, we discuss the results of Maximum-Likelihood fitting of profiles with and without a core, to the individual galaxy distributions taken from the COSMOS catalogue. In $\\S$ 4 we discuss the galaxy density distribution for composite clusters (COSMOS and ENACS), in which the individual clusters are combined.  In $\\S$ 5 we discuss the constraints provided by the outer slope of the density distributions for the parameters of the formation scenario and in $\\S$ 6 we present the conclusions. ",
        "conclusions": "We have studied the projected galaxy distributions in 77 clusters from the ESO Nearby Abell Cluster Survey. The present sample is an unbiased subset of the volume-limited ENACS sample, and thus forms a representative local (z $<$ 0.1) sample of rich (R$_{ACO}$ $>$ 1), optically selected clusters. We used both COSMOS and ENACS data to test the character of the projected galaxy distributions. In particular, we have investigated whether the galaxy distributions in rich clusters have cusps or cores in their central regions.  We have made maximum Likelihood fits to the observed distribution of COSMOS galaxies to solve for the position and the elongation of the clusters. For 15 of the 77 clusters, no reliable centre could be determined and these clusters were not considered further. Using the positions and elongations, we subsequently solved for each of the 62 remaining clusters the three parameters that describe each of the four theoretical profiles that we tested, as well as the density of background galaxies. The four model profiles that we tested against the data are the King, Hubble, NFW (Navarro, Frenk and White) and the de~Vaucouleurs profiles. Although the solutions do not converge for all of the clusters nor for all four profiles, we obtain reliable results for between 75 and 95 \\% of the clusters (depending on the model profile).  We find mean values for r$_c$, the characteristic scale of the 2-D galaxy distribution, and dispersions around the means of 128 $\\pm$ 88, 189 $\\pm$ 116, 292 $\\pm$ 191 and 1582 $\\pm$ 771 kpc, for the King, Hubble, NFW and de Vaucouleurs profiles respectively. The outer logarithmic slopes of the distributions were generalized by the usual $\\beta$-parameter, which we find to have the following average values: 1.02 $\\pm$ 0.08, 1.03 $\\pm$ 0.07, 0.61 $\\pm$ 0.05 and 7.6 $\\pm$ 0.5, for the King, Hubble, NFW and de Vaucouleurs profiles respectively, which are consistent. The average background density at the limit of the COSMOS catalogue is about 4 10$^{-5}$ galaxies arcsec$^{-2}$.  In order to investigate whether the galaxy distributions in our clusters preponderantly have cores or cusps, we have determined the likelihood ratio for the King and NFW profiles. Using all galaxies down to the COSMOS magnitude limit of about m$_{b_j}$ $\\approx$ 19.5, we find that in general the King profile is more likely to be a good representation of the data than the NFW profile. However, for the individual clusters this preference for the King profile is generally not statistically significant. If we restrict the analysis to the central regions, the significance of the preference for the King-profile fits increases, even though the number of galaxies decreases.  We have increased the statistical weight for the likelihood analysis by combining the galaxy distributions in a subset of 29 of the 62 clusters, which show a regular galaxy distribution. We take special care to avoid the creation of an artificial cusp (by taking the ellipticities into account), and to avoid the destruction of a real cusp by summing distributions with different scale lengths. We have also checked that it is unlikely that the uncertainty in the centre positions has erased a cusp. For the test we summed without scaling projected distances, after scaling with r$_c$, as well as with r$_{200}$. In all three cases we find that the King profile provides a better fit to the data than the NFW profile, at confidence levels of more than 95 \\%. Interestingly, this preference is not shared by the brighter galaxies.  Finally, we have used the outer profile slope (i.e. the result that $\\beta$ is very close to 1.0), in combination with several results from numerical models to conclude that the density parameter $\\Omega _0$ is likely to be considerably smaller than unity. In addition, the available models indicate that the Universe probably has an open geometry (i.e. no closure through $\\Lambda $ is indicated). This low implied value of $\\Omega _0$ is fully consistent with a direct determination based on the $\\frac ML$ ratios of our clusters."
    },
    "9804/astro-ph9804085_arXiv.txt": {
        "abstract": "We perform a fractal analysis of the Southern Sky Redshift Survey 2, following the methods prescribed by Pietronero and collaborators, to check their claim that the galaxy distribution is fractal at all scales, and we discuss and explain the reasons of some controversial points, through tests on both galaxy samples and simulations. We confirm that the amplitude of the two--point correlation function does not depend on the sample depth, but increases with luminosity. We find that there is no contradiction between the results of standard and non--standard statistical analysis; moreover, such results are consistent with theoretical predictions derived from standard CDM models of galaxy formation, and with the hypothesis of homogeneity at large scale ($\\sim 100$ \\h). However, for our SSRS2 volume--limited subsamples we show that the first zero--point of the autocorrelation function $\\xi(s)$ increases linearly with the sample depth,  and that its value is comparable to the radius of the maximum sphere which can be completely included in the sampled volume; this implies that the true zero--crossing point of $\\xi(r)$ has not been reached. We conclude that the apparent fractal behavior is due to a combination of a luminosity--dependent correlation amplitude and the recovering of power at larger scales in deeper samples. ",
        "introduction": "One of the pillars of the standard cosmological models is the large--scale homogeneity of the Universe (e.g. Peebles 1993). The standard statistical methods to analyze the large--scale structure, as the two--point correlation function $\\xi(r)$, are described by Peebles (1980); they rely on the definition of a mean galaxy density $\\bar{n}$, which is meaningful only if the assumption of large--scale homogeneity is true. However, at small scales the autocorrelation function of the galaxy distribution is positive and can be fitted by a power--law, which is a property of a fractal set (see Peebles 1980). On the basis of the observational evidence that galaxies are clustered in ever increasing systems, from groups and clusters to superclusters, de Vaucouleurs (1970, 1971) presented ``the case for a hierarchical cosmology''\\footnote{The idea of a hierarchical Universe has a long history, which dates back to the XVIII$^{th}$ century (see  E. Harrison, 1987, {\\em Darkness at night, A Riddle of the Universe}, Harvard University Press, Cambridge); in this century, Fournier d'Albe and Charlier were the first to propose a hierarchical (now we would say fractal) model of the Universe (see Mandelbrot 1982).}. Mandelbrot (1982) developed this concept, suggesting that the galaxy distribution in the Universe is fractal, with dimension $D = 1$.  In addition, in the last 20 years redshift surveys at increasing depths have revealed ever larger structures and voids (see e.g. Davis et al. 1982; de Lapparent et al. 1986, da Costa et al. 1994, Vettolani et al. 1997, and references therein). Einasto et al. (1986) found evidence that $r_0$ increases with sample volume, but they estimated that it should reach a value $\\sim 10$\\h~ for a fair sample of the Universe. Pietronero (1987) stressed that, if homogeneity is not reached, the correlation length $r_0$ cannot be taken as a measure of the clustering amplitude, and suggested a slightly but significantly different definition of the autocorrelation function. Adopting this approach, Coleman et al. (1988) reanalyzed the CfA1 redshift survey (Huchra et al. 1983), concluding that the distribution of galaxies was fractal to at least 20 \\h. On the other hand, Davis et al. (1988) found that $r_0$ increased as $r_0^{0.5}$, and not linearly with the depth as predicted for a simple fractal (see also Maurogordato et al. 1992). Others advocated the need for a multifractal approach (e.g. Balian \\& Schaeffer 1989; Martinez \\& Jones 1990; Martinez et al. 1990); this is however another issue and we will not discuss it: see for example the review by Borgani (1995) and references therein.  The first redshift surveys sampled relatively small volumes, and the reality and nature of correlation amplitude variations with depth remained an open --and much debated-- question. Could these variations simply reflect fluctuations due to local structures? Were they a consequence of the fractal distribution of galaxies? Or were they an indirect consequence of the dependence of clustering on galaxy luminosity?  It should be pointed out that even at a scale $\\sim 1000$ \\h~ we do not expect a {\\em perfect} homogeneity, as COBE has found evidence of anisotropy in the CMB radiation at a level $\\delta T / T \\sim 10^{-5}$ (Smoot et al. 1992), a value necessary and sufficient, at least in some standard cosmological models, to explain the formation of the observed structures. The existence of very large structures in the Universe implies that the galaxy or cluster autocorrelation function must be positive to a scale corresponding to the size of these structures; but this is not necessarily inconsistent with the measured value of the galaxy autocorrelation length $r_0 \\sim 5-6$ \\h~ for $L \\sim L_*$ galaxies (as claimed by Pietronero and collaborators), as larger structures have also a lower contrast, and the correlation of luminous matter may be significantly amplified relatively to the underlying dark matter correlation function (Kaiser 1984; Bardeen et al. 1986). Therefore, we expect that the galaxy and cluster distribution will not be perfectly homogeneous even at large scales. The claim for a fractal Universe is obviously much more stronger than that: it implies that there is no convergence to homogeneity and that it is not possible to define a mean density $\\bar{n}$ of the Universe.  In the last years Pietronero and his collaborators applied statistical indicators which do not imply a universal mean density on an ever increasing number of catalogs (see e.g. Di Nella et al. 1996, Sylos Labini, Montuori \\& Pietronero 1996, Montuori et al. 1997), claiming evidence for a galaxy fractal distribution at all scales. Their results are impressive, as recently stressed by Coles (1998), but they appear to be at variance with the results of other groups who analyzed the same catalogs with standard indicators. It is clear that the situation is unsatisfactory, despite many articles and even public debates on the subject (see Pietronero et al. 1997; Davis 1997; Guzzo 1997); while the strongest support to large--scale homogeneity comes indirectly from the high level of isotropy of the CMB radiation and from two--dimensional catalogs (Peebles 1996), there is still confusion on the interpretation of the available three--dimensional data, mainly due to the difference in the statistical indicators, which do not allow a direct and quantitative comparison of the results.  Therefore we analyzed the Southern Sky Redshift Survey 2 (SSRS2; da Costa et al. 1994) following the approach of Pietronero and collaborators, in order to independently check their claims and to answer to their criticisms (Sylos Labini et al. 1997, hereafter SMP) about the work of Benoist et al. (1996, hereafter Paper I). In section 2 we discuss the apparent dependence of $r_0$ on sample depth; in section 3 we describe the different statistics, and their relations; in section 4 we present and discuss our fractal analysis of the SSRS2 in comparison with the standard analysis; in section 5 we show that theoretical predictions derived from the standard CDM model can reproduce our results, and are therefore consistent with a homogeneous Universe; our conclusions are in section 6. ",
        "conclusions": "In this paper, we have carefully examined the claim that the Universe is a fractal, analyzing the SSRS2 and using the same statistical approach of Pietronero and collaborators.  Here are our main results:  \\begin{itemize}  \\item the SSRS2 has a luminosity dependent correlation amplitude;  \\item results obtained with the standard and fractal approach give consistent results;  \\item concerning large--scale homogeneity, observational evidence is --for the moment-- consistent with theoretical predictions derived from standard CDM models of galaxy formation, and does not require a fractal Universe;  \\item on the other hand, the first zero--point of the correlation function we measure in SSRS2 volume--limited subsamples increases linearly with the sample depth, and has approximately the same value of the radius of the largest sphere which can be included in the sample. This implies that the zero--point of $\\xi(s)$ is beyond 40 \\h, a lower limit set by very luminous galaxies (see Benoist et al. 1996; Cappi et al. 1998) and consistent with results from clusters (e.g. Cappi \\& Maurogordato 1992); on this point, we agree with SMP that measurements of $\\xi(s)$ beyond a scale corresponding to the radius of the largest sphere included in the sample are not reliable.  \\end{itemize}  We conclude that the dependence of the galaxy correlation function on the galaxy luminosity and the power still present at large scales ($\\ge 40$ \\h) can satisfactorily explain those effects interpreted by SMP as evidence of an inhomogeneous, fractal Universe."
    },
    "9804/astro-ph9804170_arXiv.txt": {
        "abstract": "Observations of clusters and super clusters of galaxies have indicated that the Universe is more dominated by baryons than ever estimated in the homogeneous cosmological model for primordial nucleosynthesis.  Recent detections of possibly low deuterium abundance in Lyman-$\\alpha$ clouds along the line of sight to high red-shift quasars have raised another potential difficulty that \\he4 is overproduced in any cosmological models which satisfy the low deuterium abundance constraint. We show that the inhomogeneous cosmological model with degenerate electron-neutrino can resolve these two difficulties. ",
        "introduction": "One of the cosmological impacts of primordial nucleosynthesis is on the universal baryon mass density $\\rho_b $.  It is of significance to answer the question how much fraction of universal mass is made of ordinary matter baryons. Homogeneous Big-Bang model for primordial nucleosynthesis~\\cite{wagoner67,copi95}, assuming the standard model for light neutrino families, predicts small $\\Omega_b$, $0.03 \\le \\Omega_b h_{50}^2 \\le 0.06$, where $\\Omega_b = \\rho_b / \\rho_c$, $\\rho_c$ is the critical density which marginally closes the Universe, and $h_{50}$ is the Hubble constant $H_0$ divided by $50\\,km/s/Mpc$. However, recent observations of rich clusters and super clusters of galaxies have indicated much larger baryon fraction, 0.1 $\\le \\Omega_b h_{50}^{3/2} \\le$ 0.3~\\cite{white93,bahcall95}.  Inhomogeneous Big-Bang model~\\cite{appligate85}--\\cite{mathews96}, which allows inhomogeneous baryon density distribution due to various physical processes in the early Universe, has been proposed in order to resolve this discrepancy. In this model difference in diffusion effects between neutrons and charged nuclei plays an important role in fluctuating density distribution to suppress overproduction of \\he4, and resultant $\\Omega_b$ is relaxed to $\\Omega_b \\sim 0.1$~\\cite{orito97}.  However, another potential difficulty has been imposed by recent observations of deuterium absorption line in Lyman-$\\alpha$ clouds along the line of sight to high red-shift quasars~\\cite{rugers96,tytler96}. Several detections~\\cite{tytler96} among them provide too small deuterium abundance to accept concordant $\\Omega_b$ which satisfies the abundance constraints on the other light elements \\he3, \\he4 and \\li7. The observed deuterium abundance still scatters largely by one order of magnitude depending on different Lyman-$\\alpha$ systems, and there are still many error sources unclear in the analysis of abundance determination. However, if these detections are real, the abundance found there is presumed to constrain most strongly the primordial abundance because these clouds are the primitive gas which still resides in the epoch of galaxy formation and has not been processed very much in its evolutionary history.  It is the purpose of this paper to propose that the inhomogeneous cosmological model with degenerate electron-neutrino can resolve these two difficulties simultaneously within the framework of the standard model for neutrino. In the next section we first discuss neutrino properties in the early Universe which can  affect strongly the primordial nucleosynthesis. We then present the results of primordial nucleosynthesis calculated in both homogeneous and inhomogeneous cosmological models in sect. 3, and the $\\Omega_b$ problem and the problem of overproduction of \\he4 are discussed in details. Finally, in sect. 4, we summarize this paper. ",
        "conclusions": "We studied the effects of lepton asymmetry of partially degenerate electron-neutrino on the primordial nucleosynthesis. Homogeneous Big-Bang model with neutrino degeneracy parameter $\\xi_{\\nu_e}$ = 0.05 can recover the concordance between \\he4 and low deuterium abundance which was found in some Lyman-$\\alpha$ clouds along the line of sight to high red-shift quasars, but the resultant $\\Omega_b$ is less than detected in rich clusters 0.1 $\\le \\Omega_b h_{50}^{3/2} \\le$ 0.3. It was found that the inhomogeneous Big-Bang model with the same degeneracy parameter can predict $\\Omega_b$ as large as 0.22, which is in reasonable agreement with observation.  This degeneracy parameter corresponds to a small chemical potential of electron-neutrino of order 10$^{-5}$ eV.  It is desirable to detect the asymmetry of background neutrinos."
    },
    "9804/astro-ph9804036_arXiv.txt": {
        "abstract": "We present a detailed analysis of a high resolution spectrum of the damped Ly$\\alpha$ system at $z_{\\rm abs}$~=~2.8112 toward PKS~0528-250. The absorption redshift is slightly larger than the emission redshift of the quasar. We estimate the column density of H$_2$ molecules $N$(H$_2$)~$\\sim$~6$\\times$10$^{16}$~cm$^{-2}$ and the fractional abundance of H$_2$, $f$~=~5.4$\\times$10$^{-5}$. The excitation temperature derived for different transitions suggests that the kinetic temperature of the cloud is $\\sim$200~K and the density $n$~$\\sim$~1000~cm$^{-3}$. The cloud therefore has a dimension of $\\sim$1~pc along the line of sight. Since it obscures the broad-line emission region, its transverse dimension should be larger than 10~pc.\\par We obtain upper limits on the column densities of C~{\\sc i} ($<$~10$^{12.7}$~cm$^{-2}$) and CO ($<$~10$^{13.2}$~cm$^{-2}$; $N$(CO)/$N$(H~{\\sc i})~$<$~7$\\times$10$^{-9}$). We suggest that the ratio $N$(H$_2$)/$N$(C~{\\sc i}) is a useful indicator of the physical conditions in the absorber. Simple photo-ionization models assuming solar relative abundances show that radiation fields with spectra similar to typical AGNs or starbursts are unable to reproduce all the constraints and in particular the surprisingly small $N$(C~{\\sc i})/$N$(H$_2$) and $N$(Mg~{\\sc i})/$N$(H$_2$) ratios. In view of the models we explored, the most likely ionizing spectrum is a composite of a UV-\"big bump\" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays. Dust is needed to explain the production of molecules in the cloud. The amount of dust is broadly consistent with the [Cr/Zn] abundance determination. ",
        "introduction": "\\label{intr}  QSO absorption line systems probe the baryonic matter over most of the history of the Universe (0~$<$~$z$~$\\la$~5).  The so-called damped Ly$\\alpha$ (hereafter DLA) systems are characterized by a very large H~{\\sc i} column density ($N$(H~{\\sc i})~$\\ga$~2$\\times$10$^{20}$ ~cm$^{-2}$), similar to the one usually seen through local spiral disks. The case for these systems to be produced by proto-galactic disks is supported by the fact that the cosmological density of gas associated with these systems is of the same order of magnitude as the cosmological density of stars at present epochs (Wolfe 1996).  Moreover the presence of heavy  elements ($Z \\sim 0.1 ~ Z_\\odot$) and the redshift evolution of metallicity suggest the ongoing star formation activities in these systems (Pettini et al.  1997), while strong metal line systems have been demonstrated to be associated with galaxies at low and intermediate $z$ (e.g.  Bergeron \\& Boiss\\'e 1991). It has also been shown that the profiles of the lines arising in the neutral gas show evidence for rotation (Wolfe 1996, Prochaska \\& Wolfe 1997).  Whether these arguments are enough to demonstrate that DLA systems arise in large disks is a matter of debate however.  Indeed simulations have shown that the progenitors of present day disks of galaxies could look like an aggregate of well separated dense clumps at high redshift. The kinematics could be explained by relative motions of the clumps with very little rotation (Haehnelt et al. 1997, Ledoux et al. 1998). Moreover, using {\\sl HST} high spatial resolution images of the field of seven quasars whose spectra contain DLA lines at intermediate redshifts (0.4~$\\la$~$z$~$\\la$~1), Le~Brun et al. (1996) show that, in all cases, at least one galaxy candidate is present within 4~arcsec from the quasar. There is no dominant morphological type in their sample: three candidates are spiral galaxies, three are compact objects and two are amorphous low surface brightness galaxies. Therefore, although the nature of the DLA systems is unclear they trace the densest regions of the Universe where star formation occurs.\\par It is thus surprising that despite intensive searches, the amount of H$_2$ molecules seems quite low in DLA systems in contrast to what is observed in our own galaxy. Two detections of H$_2$ molecules in high redshift DLA systems have been reported.  Recently Ge \\& Bechtold (1997) have found strong absorptions in the $z_{\\rm abs}$~=~1.9731 DLA system toward Q~0013--004. They derive $N$(H$_2$)~=~6.9$\\times$ 10$^{19}$~cm$^{-2}$, $b$~=~15~km~s$^{-1}$, $T_{\\rm ex}$~$\\sim$~70~K and $n$(H)~$\\sim$~300~cm$^{-3}$ for a total hydrogen column density $N$(H)~=~6.4$\\times$10$^{20}$~cm$^{-2}$. This system has by far the largest H$_2$ abundance $f$~=~2$N$(H$_2$)/[2$N$(H$_2$)~+~$N$(H~{\\sc i})] $\\sim$~0.22$\\pm$0.05 observed in high z DLA systems. However the exact number should be confirmed using a higher resolution data. Other searches have led to much smaller values or upper limits ($f$~$<$~10$^{-6}$, Black et al. 1987, Chaffee et al. 1988, Levshakov et al. 1992). Levshakov \\& Varshalovich (1985) suggested that H$_2$ molecules could be present in the $z_{\\rm abs}$~=~2.8112 system toward PKS~0528--250. This claim has been confirmed by Foltz et al. (1988) using a 1~\\AA~ resolution spectrum. The latter authors derive $N$(H$_2$)~=~10$^{18}$~cm$^{-2}$, $b$~=~5~km~s$^{-1}$, $T_{\\rm ex}$~=~100~K and log~$N$(H~{\\sc i})~=~21.1$\\pm$0.3. By fitting the damped absorption together with the Ly$\\alpha$ emission from the quasar, M\\o ller \\& Warren (1993) find log~$N$(H~{\\sc i})~=~21.35. Three Ly$\\alpha$ emission-line objects have been detected within 100$h^{-1}$~kpc from the quasar by M\\o ller \\& Warren (1993) and confirmed by Warren \\& M\\o ller (1996) to have redshifts within 200~km~s$^{-1}$ from the redshift of the DLA system ($z_{\\rm abs}$~=~2.8112 as measured on the Ni~{\\sc ii} lines by Meyer \\& York 1987). The widths of the Ly$\\alpha$ emission lines are very large ($>$~600~km~s$^{-1}$) and continuum emission could be present (Warren \\& M\\o ller 1996); this suggests that the gas is not predominantly ionized by the quasar and that star-formation may occur in the clouds, a conclusion reached as well by Ge et al. (1997). The proximity of the quasar makes the case difficult however and careful analysis is needed.\\par In this paper we use much higher spectral resolution data to reanalyze the molecular lines in this system. We present the observations in Section~2, the results  in Section~3 and discuss the low $N$(C~{\\sc i})/$N$(H$_2$) ratio inferred in Section~4. ",
        "conclusions": "\\label{s4} By fitting the different H$_2$ transitions that we detect in a high resolution spectrum of PKS~0528--250, we derive log~$N$(H$_2$)~$\\sim$~16.78 and $T_{\\rm ex}$~$\\sim$~200~K that is most certainly the true kinetic temperature of the gas. For this temperature, the relative populations of rotational levels 0 to 4 indicate that the density is of the order of 1000~cm$^{-3}$. Therefore the dimension of the neutral cloud along the line of sight is less than 1~pc. It must be noticed that the damped absorber must cover the broad line region. Indeed there is no residual flux in the bottom of the Ly$\\alpha$ absorption line which completely absorbs the Ly$\\alpha$ emission from the quasar over more than 5000~km~s$^{-1}$ (M\\o ller \\& Warren 1993). Since the dimension of the BLR in QSOs can be approximated as $R$~$\\sim$~0.3~$L_{46}^{0.5}$ where $L_{46}$ is the bolometric luminosity in units of 10$^{46}$~erg~s$^{-1}$ (Collin, private communication). The radius of the BLR in PKS~0528--250 is thus of the order of 10~pc. The transverse dimension of the damped cloud must be thus larger than 10~pc and the cloud must be quite flat.  We can derive an upper limit on the transverse dimension of the cloud by interpreting the non-dection of redshifted 21 cm absorption (Carilli et al. 1996) as an effect of partial covering factor of the continuum radio emission by the cloud. Indeed the size of the radio source is 1~arcsec or 5$h^{-1}_{75}$~kpc. If the spin temperature is equal to the kinetic temperature we derived in sect. 3.2, this implies that the covering factor should be less than 0.3 and thus the radius of the cloud  along the transverse direction is less than 1~kpc.   We determine upper limits for the C~{\\sc i}, Mg~{\\sc i} and CO column densities (log~$N$(C~{\\sc i})~$<$~12.7, log~$N$(Mg~{\\sc i})~$<$~12.8 and log~$N$(CO)~$<$~13.2). Based on simple photoionization models we conclude that (i) no simple model can reproduce at the same time the low $N$(C~{\\sc i})/$N$(H~{\\sc i}), $N$(Mg~{\\sc i})/$N$(H~{\\sc i}) ratios and the presence of molecules at the level observed; (ii) steep spectra are required to reproduce the low $N$(C~{\\sc i})/$N$(H~{\\sc i}) and $N$(Mg~{\\sc i})/$N$(H~{\\sc i}) ratios but they predict temperatures smaller than 100~K in conflict with the excitation temperature derived from the H$_2$ transitions; (iii) the only way to keep the temperature as high as 200~K is to allow for some X-ray flux heating the gas; (iv) in the framework of these models, dust is needed to produce the observed amount of molecules; (v) the gas phase abundances of carbon and especially magnesium should be smaller than 0.1 of solar.  In view of the models we explored, the most likely ionizing spectrum is a composite of a UV-\"big bump\" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays."
    },
    "9804/astro-ph9804292_arXiv.txt": {
        "abstract": "The effects that large scale fluctuations had on small scale isothermal modes at the epoch of recombination are analysed. We find that: (a)~Albeit the fact that primordial fluctuations were at this epoch still well in the linear regime, a significant nonlinear radiation hydrodynamic interaction could have taken place. (b)~Short wavelength isothermal fluctuations are unstable. Their growth rate is exponential in the amplitude of the large scale fluctuations and is therefore very sensitive to the initial conditions.  (c)~The observed CMBR fluctuations are of order the limit above which the effect should be significant.  Thus, according to their exact value, the effect may be negligible or lead to structure formation out of isothermal fluctuations within the period of recombination. (d)~If the cosmological parameters are within the prescribed regime, the effect should be detectable through induced deviations in the Planck spectrum. (e)~The sensitivity of the effect to the initial conditions provides a tool to set limits on various cosmological parameters with emphasis on the type and amplitude of the primordial fluctuation spectrum. (f)~Under proper conditions, the effect may be responsible for the formation of sub-globular cluster sized objects at particularly high red shifts. (g)~Under certain circumstances, it can also affect horizon sized large scale structure. ",
        "introduction": "The problem of structure formation in the Universe has probably been one of the foremost and most studied question in cosmology. Perhaps the greatest achievement of cosmology was the prediction and following discovery of the Cosmic Microwave Background Radiation -- the CMBR (Penzias \\& Wilson, 1965) and its anisotropy (Smoot et al. 1992). Not only did it provide support for the Big Bang theory, it proved that structure is a result of small fluctuations growing into large inhomogeneities and not vice versa. The fact that the observed fluctuations in the CMBR are small ($\\Delta T/T \\sim 10^{-5}$), naively implies that one can treat the fluctuation within the linear approximation when modeling the evolution of structure before the time of last photon scattering, i.e., one can assume that fluctuation modes of different wavelengths are completely decoupled from each other before radiation-matter decoupling.  The first to develop a linear theory for the perturbed Fridmann - Robertson - Walker metric was Lifshitz (1946), and it can be found in various text books such as Weinberg (1972), Peebles (1980) and Kolb \\& Turner (1990). The theory has since then been applied to describe cosmologies with various components and various initial parameters. Baryonic matter, radiation and other massless particles, cold (massive) or hot (light) dark matter are some of the ingredients that enter the primordial soup. While other parameters such as the Hubble constant $H_0$, the vacuum energy through the cosmological constant $\\Lambda$, and the initial spectrum affect the evolution. A review of various current cosmological models, their evolution and implication can for example be found in White et al. (1994) with emphasis on the microwave background radiation and in Primack (1997). Less current reviews that cover the basic principles of structure formation are found in the aforementioned textbooks.  One of the major parameters that affect the qualitative behaviour of our universe and is highly relevant to the present paper, is the type and form of the primordial spectrum of fluctuations.  The fluctuations are generally classified into curvature (or adiabatic) and isocurvature (or isothermal) fluctuations.  The first arise naturally in inflationary scenarios (e.g., Liddle and Lyth, 1993, and ref. therein). They are fluctuations of both the baryonic fluid and the radiation and they propagate at the adiabatic speed of sound that is equal to $c/\\sqrt{3}$ if the radiation energy density dominates. These waves are found to decay on scales smaller than the Silk scale, namely, scales comparable to galaxy sized objects (Silk, 1967). Therefore, a top-down structure formation is a natural consequence of adiabatic perturbations. Small scale objects can form after recombination (and initiate a bottom-up scenario) only if one adds the undamped perturbation of a cold component as is the case in the standard CDM model (Peebles 1982, Blumenthal et al. 1984 and Davis et al. 1985) and its variants.  Isocurvature (or isothermal) fluctuations on the other hand, are a natural consequence of topological defects formed in phase transitions (e.g, strings, monopoles and textures) or if more than one field contributes significantly to the energy density during inflation. They correspond to fluctuations that alter the entropy density but not the energy density. Unlike the first type of fluctuations, these do not suffer from Silk damping. Consequently, objects with a mass as small the post-recombination Jeans scale, namely the size of globular clusters, can form after decoupling.  It is interesting and important to note that even if the primordial spectrum was purely adiabatic, isothermal perturbations of a given wavelength are formed as the second order effect of Purcell clustering from adiabatic waves of similar wavelengths (Press \\& Vishniac, 1979). Through this effect, noninteracting particles (baryons) that are viscously coupled to a stochastically oscillating background gas (radiation fluid) undergo secular clustering. Originally, it was thought that the effect can produce a bottom-up scenario even from a pure baryon model with an adiabatic spectrum. However, the typical post-recombination Jeans scale amplitude of the isothermal waves will only be $\\rho_{iso}/\\rho\\sim (10^{-2} - 1) \\times \\left( \\delta \\rho_{ad} / \\rho\\right)^2$, with $\\delta \\rho_{ad} /\\rho$ the typical adiabatic fluctuations at horizon crossing. Namely, if the primordial spectrum is flat and adiabatic, the typical isothermal fluctuation at recombination is roughly $\\delta \\rho_{iso} /\\rho \\sim 10^{-10}-10^{-8}$. For a flat isothermal primordial spectrum, one should expect typical amplitudes of $\\rho_{iso}/\\rho \\sim 10^{-5} - 10^{-4}$.  The smallness of nonlinear effects such as the Purcell clustering (or shock waves which are of an even higher order) led to the consensus that the evolution of the fluctuations can be treated linearly and modes of different wavelengths are decoupled from each other.  This delays the nonlinear treatment to the time when radiation does not play any dynamic role anymore. At face value, it certainly appears to be the case as $\\delta\\rho /\\rho$, $v/c$, and $\\delta T/T$ are all much smaller than unity.  One should nevertheless be extremely careful when assuming linearity, especially in view of the fact that not all of the dimensionless parameters are actually smaller than unity.  One of the dimensionless numbers that appears in the solution of the nonlinear fluctuations' equations of motion and that is not small at all is the root of the radiation to gas pressure ratio.  Just before recombination one finds that $\\sqrt{P_{rad}/P_{gas}}\\approx{10}^{5}$! Consequently, we expect that the interaction between the radiation and matter at this period will have profound impact on the evolution of the fluctuations.   In this paper we examine the linear hypothesis and its validity by adding the force large scale perturbations exert on short wavelength waves. We begin in \\S2 by overviewing the problem of solving the nonlinear equations of motion and estimating the effect with a very simple analysis.  We proceed in \\S3 to write the Newtonian equations describing the evolution of short wavelength isothermal modes. In \\S4 we analyse the simplified solution, while in \\S5 we proceed to estimate the effect in a few cosmological scenarios. In \\S6 \\& \\S7, we study the possible ramifications to structure formation and study the possibility of measuring and using this effect in the study of cosmological parameters. In \\S8 we show that the effect can influence large scale structure as well. ",
        "conclusions": "We have analysed the nonlinear growth of small scale isothermal fluctuation through the induced effect that large-scale perturbations have. The equations are solved through the separation of the large scale inhomogeneities into smaller regions where they can be considered uniform, and the shear $\\Delta v_0$ exerted by the large scales is constant in space.  This can be performed because the waves that interest us are waves with very small wavelengths and a very small propagation speed, such that over the integration time, the waves practically remain in the same region of space.   Small scale isothermal waves are found to be unstable if the large scale shear velocity between the ``photon fluid'' and the baryon fluid is larger than roughly twice the isothermal speed of sound at the time of recombination. A period when the logarithmic derivative of the opacity with respect to the density does not vanish but is $-1/2$ instead.  When the shear is smaller than the critical value, the solution found is qualitatively the same as the solution for the completely decoupled equations. The shear in this case may induce some quantitative corrections but the behaviour of the system does not change.  The behaviour of the system is however radically different if the shear is larger than a critical value.  For wavelength shorter than the roughly the Jeans scale, the growth rate increases with $k$ and it saturates when $\\tau_{e\\gamma}=\\tau_o$, i.e., at a scale which is a few times smaller than the Jeans scale. The maximal growth rate attained is: \\begin{equation} r_{\\infty} = {\\alpha \\over 2} {\\Delta v_0 \\over v_s} {1\\over \\tau_e{\\gamma}}. \\end{equation} If $r_{\\infty}$ is much larger than all other rates in the system, it can be easily integrated over the duration of decoupling $\\Delta t$. The growth factor obtained is then: \\begin{equation} {\\delta_1\\over\\delta_0} \\approx \\exp (G_{\\infty}) = \\exp \\left(\\left\\langle {\\alpha \\over 2} {\\Delta v_0 \\over v_s} {1\\over \\tau_e{\\gamma}}\\right\\rangle_t \\Delta t\\right), \\end{equation} with the brackets denoting a temporal average. The isothermal waves are unstable because their sound speed is very small. Consequently, the acoustic energy in a wave with a given amplitude is very small and the amount of work needed to increase its amplitude isn't large, in fact, it is so small that nonlinear effects start to take place already at $\\delta\\rho/\\rho\\sim 10^{-5}~{\\rm to}~10^{-4}$!  Even more surprising is the fact the the typical shears found in typical cosmological scenarios is {\\em exactly} at the verge of having nonlinear effects take place! If the amplitude of the large scale fluctuations would have been an order of magnitude larger, then any isothermal perturbations, however small, could have been amplified right into the nonlinear regime. On the other hand, if the typical amplitude would have been a few times smaller, the effect would have been completely meaningless.  For typical cosmological scenarios normalized to COBE, the growth parameter is roughly $G_{rms} \\sim 1 - 5$, e.g., for an $\\Omega=1,~\\Omega_b=0.05,~h=.66$ CDM model one finds $G_{rms} \\approx 5$, if the spectral index is tilted to $n=0.9$, one finds $G_{rms}\\approx 4$, if a hot component is added at a 20\\% level to the untilted model, it falls to $G_{rms} \\sim 1$.  The fate of the regions that do reach nonlinearity is still an open question. Do these regions collapse and form black holes?  Do they form an early generation of massive stars? Do they release enough energy to the environment and affect it, or perhaps, even form structure on much larger scales? Although some speculations of what might occur do exist in the context of early small scale structure formation, it is not entirely clear. Moreover, nonlinear hydrodynamic simulations is probably unavoidable if we are to really solve the problem. The only relatively certain consequence is that the dissipation taking place after recombination heats the matter and it is likely to ionise it, raise its temperature to the ionisation temperature and leave an imprint on the CMBR in the form of a deviation from a Planckian spectrum. In scenarios where $G_{rms}\\gtrsim 5$, isothermal perturbations with an amplitude of as low as $\\delta\\rho_{iso}/\\rho\\sim10^{-6}$ will be in fact theoretically detectable in the future. An amplitude of $\\delta\\rho_{iso}/\\rho\\sim10^{-4}$ will leave 1\\% of the background sky with a detectable deviation. This fraction can change considerably if the density perturbations are not Gaussian.  The detection of deviations from a Planckian spectrum or the placing of an upper limit for them can result with interesting implications. Any deviations found will first of all directly prove the existence of primordial isothermal fluctuations. Second, such a detection will place stringent limits on cosmological parameters, as only those scenarios that produce a large enough $G_{rms}$ are capable of producing nonlinear structure, and because the amount of nonlinearity is extremely sensitive to $G_{rms}$, it is also sensitive to the exact cosmological parameters. Third, if the cosmological parameters are known with an large accuracy (for example, through the fitting of the CMBR spectrum) then the primordial isothermal perturbation spectrum at very large $k$'s can be estimated.  Even if no detection of a deviation from a Planckian spectrum is found, one can place limits on cosmological parameters. Moreover, if in the future it will be found that the cosmological parameters actually correspond to a high $G_{rms}$ model, one will be able through the lack of detection of a $y$-parameter to place extremely powerful limits on the amplitude of the isothermal component of the primordial spectrum. Through the Press \\& Vishniac effect, limits can also be placed on the adiabatic spectrum.  Another interesting implication is the possibility of transferring enough energy from the large to the small scales and considerably change the amplitude of the large-scales. This possibility relies on whether a significant volume fraction can be amplified out of the linear regime and on the maximum dissipation rate of nonlinear isothermal waves. Under favourable circumstances, the large scale amplitude can be significantly reduced such that $G_{rms}$ calculated from the observed large scale fluctuations would only be of order unity. The fact that the observed fluctuations in the CMBR correspond to values of this order raises a very interesting question.  Why do the values of $G_{rms}$ corresponding to the observed fluctuations happen to fall in a small region around unity? Is it because the universe was created with a fluctuation spectrum corresponding to this region, or, is it because $G_{rms}$ that evolved from the initial cosmological parameters was actually larger but it was naturally reduced to the oberved value?  One can summarise the several plausible scenarios according to both the value of $G_{rms}$ predicted by the cosmological model and the amplitude $\\delta = \\delta \\rho_{iso}/\\rho \\approx 10^{-10} - 10^{-5} $ of the sub Jeans scale isothermal fluctuations before the advent of recombination, these possibilities are:  \\begin{enumerate}  \\item If $G_{rms}$ is of order unity or less, the typical growth of isothermal waves is at most a few $e$-folds. The effect will be insignificant as the predicted micro degree size fluctuations are neither fluctuations of the temperature nor on a scale measurable in the near future. Moreover, there are no implications at all to structure formation.  \\item For a value of $G_{rms} \\gtrsim - (\\ln \\delta)/3 \\approx 4 - 8$, the effect will be measurable as patches in the CMBR with a distorted Planckian spectrum. As long as $G_{rms} \\lesssim - \\ln \\delta \\approx 10-20$, only a small fraction of the universe would have reached nonlinearity and formed small scale structure by the end of recombination. Under certain circumstances however, it can lead to large scale structure formation through explosive amplification.  \\item For a value of $G_{rms}$ such that $ G_{rms} \\gtrsim - \\ln \\delta \\approx 10-20$, a large fraction of the universe reaches nonlinearity by the end of recombination and small scale structure is subsequently formed in most of the universe soon after recombination. The large scale spectrum is damped by transferring a significant amount of energy to the small scales, thus, the $G_{rms}$ measured from the damped spectrum is smaller than the $G_{rms}$ calculated when neglecting this process. In some cases, $G_{rms}$ will be reduced to a values of order unity and it will mimic values on the boundary between the first and second scenarios. Note that it will not affect fluctuations in components that decouple from the matter-radiation fluid before recombination (e.g. dark matter fluctuation).  \\end{enumerate}  The analysis presented here is the first step in the investigation of the radiation-matter interaction instability. More accurate integration is needed to improve the actual transfer function for isothermal waves.  The analysis here was restricted to the evaluation of $G_{rms}$ and it does not include the actual rates which are $k$ dependent. This approximation overestimates the rate of growth of finite sized $k$'s.  In the present paper we have simulated the freezing of the instability due the freezing of recombination by stopping the integration abruptly at a given $z$. However, the equilibrium conditions and therefore the switching off of the effect depend on the isothermal wavelength as well.  By assuming this assumption, we have actually underestimated the contribution from waves of order the Jeans size.  Many of the possible implications depend on the quantitative behaviour of of the nonlinearities once they are reached. A numerical hydrodynamic study of large amplitude waves will certainly help us understand of the fate of the nonlinear objects and the possible implications they have on large scale structure as well."
    },
    "9804/astro-ph9804017_arXiv.txt": {
        "abstract": "I summarize some recent models and ideas for the formation of axisymmetrical structures of planetary nebulae and the three rings of SN 1987A, as follows. (a) I review the general role of binary companions, including brown dwarfs and planets. (b) I propose a mechanism for axisymmetrical mass loss on the AGB that may account for the axially symmetric structures of elliptical planetary nebulae and that operates for slowly rotating AGB stars, $10^{-4} \\Omega _{\\rm Kep} \\lae \\Omega \\lae 10^{-2} \\Omega_{\\rm Kep}$, where $\\Omega_{\\rm Kep}$ is the equatorial Keplerian angular velocity. (c) I propose a model for the formation of the two outer rings of SN 1987A, which is based on the numerical simulation of Soker (1989), and discuss a mechanism for their displacement from the exploding star. ",
        "introduction": "Scanning through recent images of SN 1987A (e.g. Burrows {\\it et al.} 1995) and through catalogs of planetary nebulae (PNs; e.g., Acker {\\it et al.} 1992; Schwarz, Corradi, \\& Melnick 1992; Manchado {\\it et al.} 1996) we find that the circumstellar media of many stars at their final nuclear burning phase have axisymmetrical, rather than spherical, structures. Axisymmetrical PNs which have two lobes with a morphological ``waist'' between them are termed ``bipolar PNs'' (also ``butterfly'' or ``bilobal''), while PNs which have a more elliptical than bilobal structure are termed elliptical PNs (Schwarz, Corradi, \\& Stanghellini 1993). The axisymmetrical structures of most PNs led to a debate on whether elliptical PNs can be formed through single-stellar evolution, or whether a binary companion is necessary (e.g., Fabian \\& Hansen 1979;  Livio 1982, 1998; Livio, Salzman, \\& Shaviv 1979; Webbink 1979; Morris 1981; Zuckerman \\& Gatley 1988; Pascoli 1992; Iben \\& Livio 1993; Soker 1997, 1998a; Balick {\\it et al.} 1994; Pottasch 1995; % Pollacco \\& Bell 1997; Corradi {\\it et al.} 1996; Kastner {\\it et al.} 1996). In the last decade this debate was extended to the formation of the nonspherical explosion and three rings of SN 1987A. In many PNs, as well as in the three rings of SN 1987A, there are displacements of the nebulae from the central stars, which hint at the interaction of the progenitors with wide binary companions, with close binaries having eccentric orbits, or with the ISM ($\\S 5$).  In a recent paper (Soker 1997) I suggest that four main evolutionary routes determine the degree of asymmetry of the axially symmetric structures of PNs. I then classify 458 PNs according to the process which caused their progenitors to blow axisymmetrical winds. The classification is based primarily on the morphologies of the different PNs, assuming that binary companions, stellar or substellar, are necessary for axisymmetrical mass loss on the AGB. The four evolutionary classes, according to the binary-model hypothesis, are: \\newline (a) Progenitors of planetary nebula which did not interact with any companion, and therefore they rotate extremely slowly when reaching the AGB. These amount to $\\sim 10 \\%$ of all planetary nebulae. \\newline (b) Progenitors which interact with stellar companions which avoided a common envelope, $11 ^{+ 2}_{-3} \\%$ of all nebulae. These form bipolar PNs, as is the case in symbiotic nebulae (Morris 1990; Schwarz \\& Corradi 1992; Soker 1998a). % \\newline (c) Progenitors which interact with stellar companions via a common envelope phase, $23^{+11}_{-5} \\%$ of all nebulae. These form extremely asymmetrical structures, i.e., tori, elongated elliptical PNs, and in some cases bipolar PNs. \\newline (d) Progenitors which interact with {\\it substellar} (i.e., planets and brown dwarfs) companions via a common envelope phase, $56^{+5}_{-8} \\%$ of all nebulae. These form elliptical PNs with relatively small deviation from sphericity. \\newline These numbers are compatible with other studies (e.g., Yungelson, Tutukov, \\& Livio 1993; Han, Podsiadlowski, \\& Eggleton 1995).  In $\\S 2$ I discuss the problem of angular momentum of AGB stars, which suggests that to account for the $\\sim 60 \\%$ elliptical PNs either there are many planetary systems ($\\S 3$; Soker 1996; 1997) or there is a mechanism to induce axisymmetrical mass loss from very slowly rotating AGB stars. Such a model for singly evolved very slowly rotating AGB stars is the mechanism of mode-switch to nonradial oscillations, proposed by Soker \\& Harpaz (1992). In $\\S 4$ I propose yet another model (Soker 1998c) which may operate in singly evolved AGB stars. This model is based on both magnetic activity and radiation pressure on dust. In $\\S 5$ I propose a model for the two outer rings of SN 1987A. ",
        "conclusions": ""
    },
    "9804/astro-ph9804221_arXiv.txt": {
        "abstract": "We investigate a one-zone chemophotometric evolution model of disk-disk galaxy mergers in order to clarify whether or not galaxy mergers with the widely spread  merging epoch can reproduce reasonably well the observed small scatter of the color-magnitude ($C-M$) relation in  cluster ellipticals at low and intermediate  redshift ($z<1$). We consider that merger progenitor disks begin to consume interstellar gas at moderate rate from $z \\sim 5$ and then merge to form an elliptical with the secondary starburst at $z=z_{\\rm merge}$. We find that even if the epoch of galaxy merging is rather extended ($0.3<z_{\\rm merge}<3.0$), the dispersion in the rest-frame $U-V$ color among galaxy mergers is well within the observed one ($\\sim$ 0.05 mag at $z$ = 0). We also find that the $z_{\\rm merge}$ is required to be within a certain range to keep the observed  $C-M$ relation tight at a given $z$. For example, the required range of $z_{\\rm merge}$ in galaxy mergers between Sa disks is $1.3<z_{\\rm merge}<3.0$ for cluster ellipticals at $z$ = 0.895, $0.9<z_{\\rm merge}<3.0$ for $z$ = 0.55, and $0.3<z_{\\rm merge}<3.0$  for $z$ = 0. The main reason for the derived  small scatter is  that younger stellar populations, which are formed during the secondary starburst of galaxy mergers, are formed preferentially from more metal-enriched interstellar gas. This result reinforces the Worthey's suggestion (Worthey et al. 1996) that the age-metallicity conspiracy, which means that younger stellar populations are preferentially more metal-enriched, can operate to keep the tight $C-M$ relation. These numerical results imply that the observed small scatter in the $C-M$ relation at low and intermediate redshift ($z < 1$) {\\it does not necessarily} require the coevality of elliptical galaxies in clusters or their formation at high $z$, which has been conventionally believed in the classical passive evolution picture. ",
        "introduction": "Redshift evolution of fundamental physical relations in elliptical galaxies is generally considered to give strong constraints on the formation and evolution of elliptical galaxies. For example, evolution of the color-magnitude ($C-M$) relation with redshift ($z$) suggests that elliptical galaxies are old, coeval, and homogeneous systems passively evolving after the single initial burst of star formation associated with dissipative galaxy formation at $z > 2.0$ (Arag\\'on-Salamanca et al. 1993; Ellis et al. 1997; Stanford, Eisenhardt, \\& Dickinson 1998). This classical picture of coeval elliptical galaxy formation also appears to be supported by small redshift evolution  of both the mass-to-light-ratio (van Dokkum \\& Franx 1996) and the $\\rm Mg_{2} - \\sigma$ relation (Ziegler \\& Bender 1997). The considerably tight $C-M$  relation (Bower, Lucey, \\& Ellis 1992) and the Fundamental Plane (e.g., Djorgovski \\& Davis 1987) at the present  epoch, and redshift evolution of the slope and the zero-point of the $C-M$ relation (Kodama \\& Arimoto 1997; Gladders et al. 1998; Kodama et al. 1998) furthermore seem to support the coevality of elliptical galaxy formation.  An increasing  number of recent observational results, however, shed a strong doubt on this long-standing view of elliptical galaxy formation, and suggest that there is great variety of star formation history among elliptical galaxies, such as the epoch of major star formation, the duration and efficiency of star formation (Worthey, Faber, \\& Gonzalez 1992; Faber et al. 1995; Worthey, Trager, \\& Faber 1996). In particular, Faber et al. (1995) suggested that the `apparent age spread', which is inferred from the combination of line index analysis of elliptical galaxies, amounts to $\\sim$ 10 Gyr. Schweizer \\& Seitzer (1992) found that in merger remnants with morphologically fine structures, the last merging epoch, which corresponds to elliptical galaxies formation, ranges from 4.6 Gyr to 8.0 Gyr ago. These observed  spread in  `apparent mean age' seem to be inconsistent with the aforementioned coevality of elliptical galaxy formation expected mainly from the redshift evolution of the $C-M$ relation.   The purpose of this paper is to give a plausible answer to the above apparent inconsistency in the epoch of elliptical galaxy formation. We adopt the merger scenario of elliptical galaxy formation (e.g., Toomre \\& Toomre 1972) and thereby investigate to what degree the difference in the epoch of major galaxy merging (i.e., the epoch of elliptical galaxy formation) can be allowed to preserve the observed small scatter of the $C-M$ relation of cluster ellipticals ($\\sim 0.05$ mag) at  $z$ = 0 (Bower et al. 1992), 0.55 (Ellis et al. 1997), and 0.895  (Stanford et al. 1998). We find that owing to  the age-metallicity conspiracy proposed  by Worthey et al. (1996), the observed small scatter in the $C-M$ relation can be reproduced reasonably well even in star-forming galaxy mergers with the widely spread merging epoch. This result accordingly reinforces the recent results of Kauffmann \\& Charlot (1998), in which the tight $C-M$ relation can be successfully reproduced by merger scenario of elliptical galaxy formation based on  a hierarchical clustering  scenario. This result  furthermore implies   that the previously suggested interpretation of the tightness of the $C-M$ relation at low and intermediate redshift ($z < 1$) is {\\it not unique},  thus that the formation epoch of elliptical galaxies can be more widely spread than the classical passive evolution picture predicts. Thus, the above apparent inconsistency in the interpretation of the $C-M$ relation can be due primarily  to the fact that  previous studies claiming the coevality of elliptical galaxy formation did not explore so extensively possible variety in  star formation history  of elliptical galaxies. ",
        "conclusions": "The present study predicts that even if the epoch of major galaxy merging (i.e., the epoch of elliptical galaxy formation) is rather spread, both the tightness and the slope of the $C-M$ relation can be kept owing to the age-metallicity conspiracy originally proposed by Worthey et al (1996). This result accordingly provides a heuristic explanation for the result of Kauffmann \\& Charlot (1998) in which the tight $C-M$ relation has been already reproduced in the merger scenario of elliptical galaxy formation based on the hierarchical clustering model. The conclusions derived in the present study however {\\it seem} to be inconsistent with those derived in previous ones on the redshift evolution of the slope, zero-point and tightness of the $C-M$ relation of elliptical galaxies (Bower et al. 1992; Arag\\'on-Salamanca et al. 1993; Kodama \\& Arimoto 1997). In particular, the present numerical results {\\it seem} to disagree with those of Kodama \\& Arimoto (1997) and Kodama et al. (1998) (see also Gladders et al. 1998), which claim that the considerably less significant evolution of the slope of the $C-M$ relation rejects the age spread larger than 1 Gyr among elliptical galaxies. The apparent disagreement between the present study and the previous ones (e.g., Kodama \\& Arimoto 1997; Kodama et al. 1998) is due essentially to the fact that the previous studies inevitably  have over-interpreted  the redshift evolution of the $C-M$ relation owing to the ad hoc assumption adopted in the previous studies. Although the previous studies are considerably sensible and valuable, it is important  to point out the ad hoc assumptions adopted in the previous studies and thereby clarify the reason why the present conclusions are not consistent with those derived by the previous studies of Kodama \\& Arimoto (1997) and Kodama et al. (1998). The following three are the ad hoc assumptions which inevitably lead the previous studies  to draw the strong and general conclusion that formation of elliptical galaxies (especially in the cores of clusters) are {\\it as a whole}  coeval and occurred at high redshift. First is that elliptical galaxies are formed by {\\it only one} initial starburst. Owing to this assumption, time evolution of global colors of elliptical galaxies depends exclusively on the epoch of initial burst of star formation (i.e., the epoch of elliptical galaxy formation in the previous  study). As a result of this, the age difference between elliptical galaxies (i.e., the difference of the epoch of elliptical galaxy formation in the previous studies) can be more clearly reflected on the redshift evolution of the slope of the $C-M$ relation in the previous studies. Accordingly the observed less significant evolution of the slope of the $C-M$ relation is more likely to be interpreted as an evidence that supports the coevality of elliptical galaxy formation. It is certainly reasonable to claim that the observed evolution of the $C-M$ relation reject the `pure age' sequence model which demands that less luminous ellipticals have younger age. However, it seems not to be so reasonable to draw strong and general conclusion that elliptical galaxies are formed at $z>2$ {\\it only} from the redshift evolution of the $C-M$ relation. Considering the first ad hoc assumption in the previous studies, what is more accurate and plausible interpretation on the observed evolution of the $C-M$ relation is just that the formation of {\\it stellar populations} in {\\it some} elliptical galaxies in the cores of {\\it some} clusters (not the formation of galaxies with structural and morphological properties similar to those of ellipticals) can be coeval and occurred at higher redshift ($z > 2$). The second is that an elliptical galaxy in a cluster of galaxies at higher redshift is a precursor of an elliptical galaxy in a cluster at lower redshift. The third is that a cluster of galaxies observed at higher redshift is a precursor of a cluster of galaxies at lower one. These two ad hoc assumptions actually enable us to discuss the origin of elliptical galaxies in a more general way and thus lead us to draw more general conclusions on the formation epoch of elliptical galaxies. However, since there are no observational evidences which can provide the firm physical basis for the above assumptions at least now, it is questionable to give any general conclusions on the coevality of elliptical galaxy formation. Thus, these three assumptions adopted in the previous studies inevitably lead them to provide the strong and general conclusion that formation of elliptical galaxies are coeval and occurred at higher redshift. The present study, on the other hand, does not adopt the above three ad hoc assumptions, and rather relaxes these assumptions. Furthermore the present study instead allows both continuous and moderate star formation (not strong initial starburst) and the secondary starburst associated with galaxy merging, and assumes that the epoch of morphological transformation (into ellipticals) does not necessarily coincide with the epoch of galaxy formation (i.e., the epoch when the star formation begins). The evolution of the $C-M$ relation in the present study consequently  does not depend so strongly on the difference in the formation epoch between  elliptical galaxies (i.e., the epoch of major galaxy merging with the secondary starburst). As a result of this, the present merger model predicts that even if the formation epoch of elliptical galaxies (i.e., the epoch of galaxy merging) are rather spread, both the slope and tightness of the $C-M$ relation can be kept. Thus, the essential reason for the aforementioned  apparent disagreement on the coevality of elliptical galaxy formation is that the present study does not adopt the above three ad hoc assumptions whereas the previous studies do. The interpretation on the redshift evolution of the $C-M$ relation in each model can depend strongly on the assumptions adopted by each model. It is safe for us to say that it is not clear, at least now, which of the two different conclusions on the coevality of elliptical galaxy formation is more plausible and reasonable. However, considering the above three ad hoc assumptions adopted in the previous studies, what is more reasonable interpretation on the redshift evolution of the $C-M$ slope is that only {\\it stellar populations} (not elliptical morphology) in {\\it some} ellipticals located in  the cores of {\\it some} clusters of galaxies are formed at higher redshift. We should not draw any {\\it general} conclusions from the redshift evolution of the slope of the $C-M$ relation.   Environmental difference of stellar populations (in particular, the existence of intermediate-age population) in early-type galaxies has been already indicated by a number of observational studies (e.g., Bower et al. 1990; Rose et al. 1994; Mobasher \\& James 1996). On the other hand, the tightness and the slope of the $C-M$ relation of early-type galaxies are observationally revealed not to depend so strongly on galaxy environments. These two apparently inconsistent observational results on spectrophotometric properties of elliptical galaxies have called into the following question: ``Why does not   the $C-M$ relation of early-type galaxies depend strongly on galaxy environments (e.g., between rich clusters and poor ones), though stellar populations and star formation histories in early-type galaxies probably depend on galaxy environments?'' To give a plausible answer for this question seems to be important because the above apparently inconsistent observational results give us valuable information both on the environmental difference in the details of physical processes of elliptical galaxy formation and on a certain mechanism for the tight $C-M$ relation. However, no extensive theoretical studies have yet addressed the above important question. The present study has shown that the age-metallicity conspiracy, which is achieved by younger and more metal-enriched stellar populations created in the secondary starburst of galaxy mergers, allows both the apparent age spread of elliptical galaxies and the tightness of the  $C-M$ relation. This result seems to  provide a clue to the above question. Since the {\\it real} question concerning the tight $C-M$ relation is not to determine the typical epoch of elliptical galaxy formation but to give a convincing explanation for the reason why possible diversity in star formation histories of elliptical galaxies can allow the tight $C-M$ relation, more extensive theoretical studies including more variety of star formation history of elliptical galaxies and its likely dependence on galaxy environments are certainly worth for our deeper understanding of the origin of the tight $C-M$ relation.   The present numerical results  are consistent with recent observational results which suggest that coeval elliptical galaxy formation with initial starburst at higher redshift ($z >  2.0$) is not promising. Kauffmann, Charlot, \\& White (1996) revealed that only about one-third of bright E/S0 galaxies in the sample of Canada-France Redshift Survey were already in the passive evolution phase at $z \\sim 1.0$. Franceschini et al. (1997) found a remarkable absence of early-type galaxies at $z > 1.3$ in the $K$-band selected sample of early-type galaxies in the Hubble Deep Field (HDF), which suggests either that early-type galaxies are formed by galaxy merging with less prominent star formation or that a dust-polluted interstellar gas obscures forming elliptical galaxies till $z = 1.3$. Zepf (1997) demonstrated that strong deficit of galaxies with extremely red colors in the HDF means that the formation epoch of typical elliptical galaxies is   $z < 5.0$. Sample galaxies in these studies are selected from {\\it field ellipticals}, which possibly have star formation histories different from those of {\\it cluster ellipticals}. Accordingly, it might not be plausible to derive strong conclusions on the formation epoch of ellipticals. However these observational results together with the present results seem to support the merger scenario which can naturally predict that the epoch of elliptical galaxy formation is rather extended ranging from high redshift to moderate one.    Thus we  have succeeded in pointing out  that even if the epoch of elliptical galaxy formation (i.e., the epoch of major disk-disk galaxy merging, in this study) is rather widely spread, the tightness of the $C-M$ relation at low and intermediate redshift can be kept reasonably well. This result suggests  that coevality of elliptical galaxy formation, which has been  conventionally believed in the classical passive evolution picture, is {\\it not unique} interpretation  for the small  scatter of the $C-M$ relation. This  furthermore implies that {\\it only} the tightness of the $C-M$ relation at a given redshift {\\it does not necessarily} give strong constraints on the formation epoch of elliptical galaxies. Worthey et al. (1996) have already pointed out that the age-metallicity conspiracy can keep both the tightness of  the Fundamental Plane and that of the $C-M$ relation in elliptical galaxies. The present numerical study, which is different from the Worthey's single stellar population analysis, has confirmed that the proposed age-metallicity conspiracy can actually operate to keep convincingly the tightness of the $C-M$ relation of ellipticals formed by  disk-disk galaxy mergers. The present chemophotometric evolution model is, however, not so elaborated and realistic in that this model neither includes continuous gas accretion/merging expected from a specific cosmology (e.g., Baugh, Cole, \\& Frenk 1996; Kauffmann \\& Charlot 1998) nor considers important dynamical effects of galaxy merging on chemical and photometric evolution of galaxies (Bekki \\& Shioya 1998). Accordingly  it is our future study to confirm that the results derived in the present preliminary study can hold even for more sophisticated and realistic merger models. Furthermore, we should  check whether or not observed redshift evolution of other fundamental relations such as the $\\rm  Mg_{2} - \\sigma$ relation (Ziegler \\& Bender 1997), the Fundamental Plane (van Dokkum \\& Franx 1996), and the abundance ratio of [Mg/Fe] can be also reproduced self-consistently by our future merger model."
    },
    "9804/hep-ph9804378_arXiv.txt": {
        "abstract": "In this paper we present the first analytic model for vorton formation. We start by deriving the microscopic string equations of motion in Witten's superconducting model, and show that in the relevant chiral limit these coincide with the ones obtained from the supersonic elastic models of Carter and Peter. We then numerically study a number of solutions of these equations of motion and thereby suggest criteria for deciding whether a given superconducting loop configuration can form a vorton. Finally, using a recently developed model for the evolution of currents in superconducting strings we conjecture, by comparison with these criteria, that string networks formed at the GUT phase transition should produce no vortons. On the other hand, a network formed at the electroweak scale can produce vortons accounting for up to $6\\%$ of the critical density. Some consequences of our results are discussed. ",
        "introduction": "\\label{v-in} As first pointed out by Witten \\cite{witten}, cosmic strings can in some circumstances (typically when the electromagnetic gauge invariance is broken inside the string) behave as `superconducting wires' carrying large currents and charges---up to the order of the string mass scale in appropriate units. The charge carriers can be either bosons or fermions (see \\cite{vs} for a review). The former type occurs when it becomes energetically favourable for a charged Higgs field to have a non-zero vacuum expectation value in the string core; the latter happens when fermions couple to the string fields creating fermion zero modes.  It is well known that arbitrarily large currents are not allowed---there is a critical value beyond which the current saturates. In other words, for large enough winding number per unit length, the superconducting condensate is quenched down, suppressing the current flow. Also, the current can decay by magnetic flux-line tunnelling; this can be used to impose constraints on allowed particle physics models.  If superconducting strings carry currents, they must also carry charges of similar magnitude. This includes not only charges trapped at formation by the Kibble mechanism but also the ones due to string inter-commuting between regions of the string network with different currents. Just like with currents, charge densities cannot have arbitrarily large magnitude---there is a limit beyond which there will no longer be an energy barrier preventing the charge carriers from leaving the string.  A rather important point is that the presence of charges on the string tends to counteract  the current quenching effect discussed above. In fact, numerical simulations of contracting string loops at fixed charge and winding number have shown \\cite{davsh} that a `chiral' state with equal charge and current densities is approached as the loop contracts. In this limiting chiral case, quenching is in fact eliminated completely. This has several important consequences. Strings that have trapped charges as a consequence of a phase transition can become superconducting even if the formation of a condensate was otherwise energetically unfavoured. More importantly, a string with both a charge and a current density will have a non-zero angular momentum.  In the cosmological context, these strings would of course interact with the cosmic plasma, originating a number of interesting consequences. The most remarkable of these, however, has to do with the evolution of string loops. If a superconducting string loop has an angular momentum, it is semi-classically conserved, and it tries to resist the loop's tension. This will at least increase the loop's lifetime. If the current is too large, charge carriers will leave the string accompanied by a burst of electromagnetic radiation, but otherwise it is possible that dynamically stable loops form. These are called vortons \\cite{vor}---they are stationary rings that do not radiate classically, and at large distances they look like point particles with quantised charge and angular momentum. Their cosmological significance comes from the fact that they provide very strong constraints on allowed particle physics models, since they behave like non-relativistic particles. According to current belief \\cite{vor,vor2}, if they are formed at high enough energy scales they are as dangerous as magnetic monopoles, producing an over-density of matter in disagreement with observations. On the other hand, low-mass vortons could be a very interesting dark matter candidate. Understanding the mechanisms behind formation and evolution is therefore an essential cosmological task.  The overwhelming majority of the work done on cosmic strings so far was concerned with the structureless Goto-Nambu strings (but see \\cite{carternew} and references therein for some exceptions). In the case of work on vortons, this means that somewhat {\\em ad-hoc} estimates had to be made for some properties of the cosmic string network---notably for microscopic quantities such as current and charge densities. This is despite the fact it has been recognised a long time ago that, even though they might be computationally very useful \\cite{ms,ms2,ms2a}, Goto-Nambu models cannot realistically be expected to account for a number of cosmologically relevant phenomena, due to the very limited number of degrees of freedom available. Two such phenomena are the build-up of small-scale structure and charge and current densities.  In this paper we fill this important gap by discussing the problem of vorton formation in the context of the superconducting string models of Witten \\cite{witten} and of Carter and Peter \\cite{wcp} (sections \\ref{v-wt} and \\ref{v-cp}). Strangely enough, the issue of the conditions for vorton formation has been so far neglected with respect to those of their stability and cosmological consequences. We will start by introducing these models and determining the microscopic string equations of motion in each case. It will be shown that in the relevant chiral limit these equations coincide---this also provides the first conclusive evidence of the validity of the supersonic elastic models of Carter and Peter \\cite{wcp}.  We then proceed to study the evolution of a number of loop solutions of these equations numerically (sections \\ref{v-fl} and \\ref{v-ex}), and from the results of this analysis parameters will be introduced which characterise the loop's ability to evolve into a vorton state (section \\ref{v-vt}). Finally, we discuss a very simple phenomenological model for the evolution of the superconducting currents on the long cosmic string network \\cite{mss}, based on the dynamics of a `superconducting correlation length' (sections \\ref{v-cr}--\\ref{v-ff}). Using this model we can therefore estimate the currents carried by string loops formed at all relevant times, and thus (in principle) decide if these can become vortons (section \\ref{v-rs}) and calculate the corresponding density (section \\ref{v-den}).  Based on our results, we don't expect any GUT vortons to form at all. This is essentially because the friction-dominated epoch is very short for GUT-scale strings \\cite{ms}, so their currents and charges are never large enough to prevent them from becoming relativistic---and therefore liable to losses. Even if they did form, they wouldn't be in conflict with the standard cosmological scenario if they decayed soon after the end of the friction-domination epoch.  Hence we conclude that, in contrast with previously existing estimates \\cite{vor,vor2}, one cannot at the moment rule out GUT superconducting string models. We should point out at the outset that there are essentially three improvements in the present work which justify the different end result for GUT-scale strings. Firstly, by analysing simple (but physically relevant) loop solutions of the microscopic string equations of motion for the Witten model, we can get a much improved idea of how superconducting loops evolve and of how (and under which conditions) they reach a vorton state. Secondly, by using a simple model for the evolution of the currents on the long strings \\cite{mss} we can accurately determine the typical currents on each string loop at the epoch of its formation. Finally, the use of the analytic formalism previously introduced by the present authors \\cite{ms,ms2a} allows us to use a quantitative description throughout the paper, and in particular to determine the loop sizes at formation.  As will become clear below, when taken together these allow a detailed analysis of the process of vorton formation to be carried out, either in the Witten model (as is done in this paper) or any other that one considers relevant. In contrast, note that Davis \\& Shellard \\cite{vor} restrict themselves to the particular case of the initial Brownian Vachaspati-Vilenkin loops with Kibble currents, and do not consider the subsequent evolution of the network. On the other hand, Brandenberger {\\em et al.} \\cite{vor2} make rather optimistic order-of-magnitude estimates about the process of relaxation into a vorton state. As it turns out, for high energy GUT scales, all these loops become relativistic before reaching a vorton state. Finally, neither of these treatments has the benefit of a quantitative model for the evolution of the long-string network \\cite{ms} which allows one to accurately describe the process of loop production.  On the other hand, as we lower the string-forming energy scale we expect more and more efficient vorton production, and the 'old' scenario still holds. Therefore intermediate-scale superconducting strings are still ruled out, since they would lead to a universe becoming matter-dominated earlier than observationally allowed. Finally, at low enough energy scales, vortons will be a dark matter candidate. For example, for a string network formed around $T\\sim10^2\\,GeV$ (typical of the electroweak phase transition) they can provide up to $6\\%$ of the critical density. A more detailed discussion of these issues is left to a forthcoming publication \\cite{inprep}.  Throughout this paper we will use fundamental units in which $\\hbar=c=k_B=Gm^2_{Pl}=1$. ",
        "conclusions": "\\label{v-cc} In this paper we have presented the first rigourous study of the cosmological evolution of superconducting strings in the limit of chiral currents. We have shown that in this limit the elastic string model of Carter \\& Peter\\cite{wcp} coincides with the model derived from first principles by Witten\\cite{witten}.  By analysing physically relevant loop solutions of the microscopic equations of motion for these strings, we have verified that the effect of frictional damping is crucial for vorton formation. We then defined suitable parameters characterising the evolution of these loops, and in particular whether or not they become vortons. In particular, we have established the usefulness of the `stability parameter' ${\\overline n}$. In general, it is more difficult to form vortons when the string-forming phase transition is of first order. This is because such networks produce, during their evolution in the stretching regime, loops with a size close to that of the horizon; these will therefore be significantly affected by expansion, which tends to decrease the fraction of the loops's energy in the current---whereas friction tends to increase it.  After introducing a simple `toy model' for the evolution of currents on the strings \\cite{mss}, we have considered the cases of first and second-order GUT-scale string-forming and superconducting phase transitions (which is the most favourable GUT case of vorton formation since frictional forces can act longer). We have presented evidence suggesting that GUT-scale string networks might well produce no vortons, and that even if they do, this will not necessarily rule out such models. This is in contradiction with previous, less detailed studies \\cite{vor,vor2}, and hence calls for a re-examination of a number of cosmological scenarios involving superconducting strings. Notably, these strings could be at the origin of the observed galactic magnetic fields \\cite{msmag}.  Finally, we have explicitly calculated the vorton density in two `extreme' cases to illustrate the method that one should follow once the microphysical properties of these networks are known in more detail. For electroweak-scale string networks, we have found that vortons can produce up to about $6\\%$ of the critical density of the universe. On the other hand, it is conceivable that superconducting string networks formed at an energy scale $T\\sim10^4-10^6\\,GeV$ (depending on details of the model) can solve the dark matter problem.  The detailed analysis presented in this paper for GUT stings can obviously be extended to other energy scales---this will be the subject of a forthcoming publication \\cite{inprep}. Obviously, as we lower the energy scale, the frictional force becomes more and more important and acts for a longer time. Hence the vorton-forming region of parameter space increases, and by the electroweak scale almost all loops chopped off the long-string network will become vortons. We therefore conclude that in addition to the low-$G\\mu$ regime (which as we saw includes the electroweak scale) where vortons can be a source of dark matter and to an intermediate-$G\\mu$ range in which vortons would be too massive to be compatible with standard cosmology (thereby excluding these models), there is also a high-$G\\mu$ regime (of which the GUT scale is part) in which vortons don't form at all and therefore no cosmological constraints based on them can be set. It is then curious (to say the least) that vorton constraints can be used to rule out cosmic string models in a wide range of energy scales $G\\mu$, but not those formed around the GUT or the electroweak scales, where cosmic strings can be cosmologically useful."
    },
    "9804/astro-ph9804235_arXiv.txt": {
        "abstract": "We investigate the possibility of accounting for the currently inferred primordial abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li by big bang nucleosynthesis in the presence of baryon density inhomogeneities plus the effects of late--decaying massive particles (X), and we explore the allowed range of baryonic fraction of the closure density $\\Omega_{b}$ in such context. We find that, depending on the parameters of this composite model (characteristic size and density contrast of the inhomogeneities; mass--density, lifetime, and effective baryon number in the decay of the X particles), values as high as $\\Omega_{b}h_{50}^{2}\\simeq 0.25-0.35$ could be compatible with the primordial abundances of the light nuclides. We include diffusion of neutrons and protons at all stages, and we consider the contribution of the X particles to the energy density, the entropy production by their decay, the possibility that the X--products could photodissociate the light nuclei produced during the previous stages of nucleosynthesis, and also the possibility that the decay products of the X--particles would include a substantial fraction of hadrons. Specific predictions for the primordial abundance of Be are made. ",
        "introduction": "Standard homogeneous big bang nucleosynthesis could have produced the observationally inferred primordial abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li, provided that the baryon fraction of the cosmic closure density $\\Omega_{b}$ would lie in the range:  $$0.04\\lapprox \\Omega_{b}h_{50}^{2}\\lapprox 0.08\\eqno(1)$$  \\noindent where $h_{50}$ is the Hubble constant in units of 50 km s$^{-1}$ Mpc$^{-1}$ (Walker et al. 1991; Copi, Schramm, \\& Turner 1995). For the long--time most favored cosmological model, a flat Universe with $\\Omega_{M} = 1$ and $\\Omega_{\\Lambda} = 0$ (those being, respectively, the fractional contributions of matter and vacuum energy densities to the closure density), the upper bound to $\\Omega_{b}$ would mean that most matter in the Universe should be in nonbaryonic form. Given the far--reaching implications of the dominance of nonbaryonic dark matter, possible alternatives to homogeneous big bang nucleosynthesis have been explored, especially during the last 15 years or so.  The suggestion that the quark--hadron phase transition might be first--order and generate baryon inhomogeneities (Witten 1984) led to the calculation of the possible effects on primordial nucleosynthesis (Applegate \\& Hogan 1985; Applegate, Hogan, \\& Scherrer 1987; Malaney \\& Fowler 1988). The goal was to see whether inhomogeneous big bang nucleosynthesis with $\\Omega_{b} = 1$ might account for the primordial light--element abundances. Besides a first--order quark--hadron phase transition, other mechanisms might also generate baryon inhomogeneities. Much of the work in this line is reviewed by Malaney \\& Mathews (1993). However, the recent studies, treating accurately the coupling between  baryon diffusion and nucleosynthesis, show that the upper limit on $\\Omega_{b}$ set by the light--element abundances does not significantly differ from that obtained for homogeneous big bang nucleosynthesis (Mathews, Schramm, \\& Meyer 1993; Thomas et al. 1994). This last conclusion, though, has very recently been challenged by Orito et al. (1997), who explore the dependence of primordial nucleosynthesis on the geometry of baryon inhomogeneities and find that cylindrical geometry might allow to satisfy the observational constraints with baryon fractions as high as $\\Omega_{b}h_{50}^{2}\\lapprox 0.2$.  A different approach has been to explore the possible modifications of the yields from homogeneous big bang nucleosynthesis by the effects of the decay of unstable massive particles ($M\\gapprox few\\ GeV$), produced at earlier stages in the evolution of the Universe and with half--lives longer than the standard nucleosynthesis epoch ($\\tau_{x}\\gapprox  10^{4}\\ s$) (Audouze, Lindley, \\& Silk 1985; Dom\\'{\\i}nguez--Tenreiro 1987; Yepes \\& Dom\\'{\\i}nguez--Tenreiro 1988; Dimopoulos et al. 1988). Gravitinos produced during reheating at the end of inflation are a possible example of such particles. In Dimopoulos et al. (1988), the emphasis is put on the resulting hadron cascade. The main problem encountered in this model is the predicted overproduction of $^{6}$Li: $^{6}Li/^{7}Li\\gg 1$, whereas observations show $^{6}Li/^{7}Li\\lapprox 0.1$.  Although they have only been considered separately, baryon inhomogeneities and the presence of unstable massive particles decaying when the Universe has already cooled down below $T_{9}\\simeq 0.4$ are by no means mutually exclusive. Here we explore their combined effects on the primordial abundances of the light elements. The parameter space now has, of course, a dimension which is the sum of those for the two separate cases: characteristic size and density contrast of the inhomogeneities, mass--density, lifetime, and mode of decay of the massive particles. We find that there are regions in such extended parameter space where values of $\\Omega_{b}$ as high as $\\Omega_{b}h_{50}^{2}\\simeq 0.35$ would still be compatible with the primordial abundances of the light nuclides inferred from observations. Such values of $\\Omega_{b}$ are of the same order as the low values for $\\Omega_{M}$ currently derived from a variety of sources, including high--redshift supernova searches (Perlmutter et al. 1998; Garnavich et al. 1998). Our results thus suggest again the possibility that all the matter in the Universe could be baryonic.  On the other hand, recent determinations of the D abundance in high--redshift QSO absorbers, when confronted with the currently inferred primordial $^{4}$He abundance, might be in conflict with the predictions of standard, homogeneous big bang nucleosynthesis for $N_{\\nu} = 3$ (Steigman 1998): the ``low'' high--redshift D abundances (which appear more reliable) would indicate too high a value of $\\Omega_{b}$ to be compatible with that corresponding to the $^{4}$He abundance. Since it is hard to tell whether this conflict points to new physics or just to systematic errors in the derivation of abundances, Steigman, Hata, \\& Felten (1998) have discarded the constraint on $\\Omega_{b}$ from standard big bang nucleosynthesis and turned to other observational constraints to determine the key cosmological parameters. The results from our composite model, by showing how minor deviations from the standard hypotheses can produce agreement with the primordial abundances inferred from observations, support that attitude. Besides, as we will see, the combined effects of inhomogeneities plus late--decaying particles might solve the conflict between D and $^{4}$He abundances. ",
        "conclusions": "We have shown, by means of a simple model, that the combined effects on big bang nucleosynthesis of baryon inhomogeneities plus the decay of unstable, relatively long--lived massive particles, giving rise to both electromagnetic and hadron cascades, might be to allow agreement with the primordial light--element abundances inferred from observations for values of $\\Omega_{b}$ much higher that those allowed by standard, homogeneous nucleosynthesis. The upper limit might be as high as $\\Omega_{b}h_{50}^{2}\\simeq 0.35$. The values obtained here are of the same order as the low $\\Omega_{M}$ values now being derived from a variety of sources and, therefore, they pose in new terms the question of whether all matter in the Universe could be baryonic. A testable prediction of the model is the production of a $^{9}$Be abundance that is of the order of current observational upper limits.  On the other hand, in the parameter region of our model where there is agreement between predicted and observationally inferred primordial light--element abundances, given values of $\\Omega_{b}$ (or, equivalently, $\\eta_{10}$) always predict ``low'' D abundances (in the sense of the high--redshift abundances referred to in the Introduction), thus potentially eliminating the conflict with the $^{4}$He abundance pointed out by Steigman (1998).  The model presented here deals with inhomogeneities in a very simplified way. A futher step will be to examine the effects of the geometry of the density fluctuations on the outcome. Orito et al. (1997) have already shown that cylindrical shell geometry alone (without the extra effects of late--decaying particles) might allow $\\Omega_{b}\\lapprox 0.2$ (but for density contrasts $R\\sim 10^{6}$, much higher than those considered here). Another extension of the model will be to consider particles with shorter half--lives, decaying at the time when thermonuclear reactions are still taking place."
    },
    "9804/astro-ph9804145_arXiv.txt": {
        "abstract": "The neutrino emissivities resulting from direct URCA processes in neutron stars are calculated in a relativistic Dirac-Hartree approach in presence of a magnetic field. In a quark or a hyperon matter environment, the emissivity due to nucleon direct URCA processes is suppressed relative to that from pure nuclear matter. In all the cases studied, the magnetic field enhances emissivity compared to the field-free cases. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804286_arXiv.txt": {
        "abstract": "We present a detailed morphological analysis of the galaxy populations in the first two clusters to be completed in an extensive observational study of nine high-redshift clusters of galaxies (Oke, Postman \\& Lubin 1988). These two clusters, CL0023+0423 and CL1604+4304, are at redshifts of $z = 0.84$ and $z = 0.90$, respectively. The morphological studies are based on high-angular resolution imagery taken with WFPC2 aboard the {\\it Hubble Space Telescope}. These data are combined with deep, ground-based $BVRI$ photometry and spectra taken with the Keck 10-meter telescopes. The morphological classifications presented in this paper consist of two parts. Firstly, we provide a quantitative description of the structural properties of $\\sim 600$ galaxies per cluster field using the Medium Deep Survey automated data reduction and object classification software (Griffiths et al.\\ 1994; Ratnatunga, Ostrander \\& Griffiths 1997). This analysis includes the galaxy position, photometry, and best-fit bulge+disk model. Secondly, for the brightest subsample of $\\sim 200$ galaxies per cluster field, we provide a more detailed morphological description through a visual classification based on the revised Hubble classification scheme (e.g.\\ Sandage 1961; Sandage \\& Bedke 1994).  Based on these classifications, we have examined the general relation between galaxy morphology and other photometric and spectral properties. We find that, as expected, the elliptical and S0 galaxies are redder, on average, than the spirals and irregulars. In addition, there is a strong correlation between morphology and spectral type. Of the galaxies that are visually classified as ellipticals, the majority show K star absorption spectra which are typical of nearby, red early-type galaxies; however, a few are actually blue compact galaxies with spectra characterized by fairly strong, narrow emission lines. Normal late-type galaxies typically have spectra with blue colors and [\\ion{O}{2}] emission, while the presence of strong star-formation features, such as extremely high equivalent width [\\ion{O}{2}], ${\\rm H\\beta}$, and/or [\\ion{O}{3}] emission, is always accompanied by peculiar morphologies which suggest recent mergers or interactions.  We have used the statistical distributions of cluster galaxy morphologies to probe the overall morphological composition of these two systems.  This analysis reveals that the two clusters contain very different galaxy populations.  CL0023+0423 has a galaxy population which is more similar to groups of galaxies and the field.  This system is almost completely dominated by spiral galaxies. CL1604+4304, however, has a morphological composition which is more typical of a normal, present-day cluster; early-type galaxies comprise $\\sim 76\\%$ of all galaxies brighter than $M_{V} = -19.0 + 5~{\\rm log}~h$ in the central $\\sim 0.5~h^{-1}~{\\rm Mpc}$. The ratio of S0 galaxies to ellipticals in this cluster is $1.7 \\pm 0.9$, consistent with local cluster populations. The morphological results support the conclusions of the dynamical analysis presented in the second paper of this series (Postman, Lubin \\& Oke 1998). CL0023+0423 consists of two galaxy groups which are separated by $\\sim 2900~{\\rm km~s^{-1}}$ in radial velocity. CL1604+4304, on the other hand, has a velocity distribution indicating that it is already well-formed and relaxed. The morphological composition, velocity dispersion, and implied mass of the CL1604+4304 system are consistent with an Abell richness class 2 or 3 cluster. ",
        "introduction": "The study of the galaxy populations of rich clusters provides important constraints on the formation mechanisms of both clusters and galaxies. Present--day clusters show a distinct correlation between the structure of the cluster and the galaxy population. Irregular, open clusters, such as Virgo, are spiral--rich. These systems show no single, central condensation, though the galaxy surface density is at least five times as great as the surrounding field ($n_{\\rm gal} > 30~h^{3}~{\\rm galaxies~{Mpc}^{-3}}$). These clusters are often highly asymmetric and have significant degrees of substructure.  Dense, centrally concentrated clusters, such as Coma, contain predominantly early--type galaxies in their cores (Abell 1958; Oemler 1974; Dressler 1980a,b; Postman \\& Geller 1984).  These clusters have a single, prominent concentration among the bright member galaxies and typically display a high--degree of spherical symmetry, though this does not preclude evidence of some substructure. Central densities can reach as high as $10^{4}~h^{3}~{\\rm galaxies~{Mpc}^{-3}}$ (e.g\\ Bahcall 1975; Dressler 1978). In these regions, spiral galaxies comprise less than 10\\% of the cluster population, while elliptical (E) and S0 galaxies make up 90\\% or more of the population. The ratio of S0s to ellipticals is typically S0/E $\\sim 2$ (Dressler 1980a). The galaxy content of clusters is part of the general morphology--density relation of galaxies; as the local density increases, the fraction of elliptical and S0 galaxies increases, while the fraction of spiral galaxies decreases (Hubble 1936; Dressler 1980a,b; Postman \\& Geller 1984).  Previous studies of clusters of galaxies at $z < 1$ have revealed significant evolution in the morphology and the color of the cluster members.  One of the most notable of these changes is the progressive blueing of cluster's galaxy population with redshift, a trend first observed by Butcher \\& Oemler (1984). They found that the fraction of blue galaxies in a cluster is an increasing function of redshift, indicating that clusters at redshifts of $z \\sim 0.5$ are significantly bluer than their low--redshift counterparts. At redshifts of $z \\sim 0.4$, the fraction of blue galaxies is $\\sim 20\\%$.  Recent HST image data reveal that many of these blue galaxies are either ``normal'' spirals or have peculiar morphologies, resulting in non--elliptical fractions which are 3 to 5 times higher than the average current epoch cluster (Dressler et al.\\ 1994; Couch et al.\\ 1994; Oemler, Dressler \\& Butcher 1997; Dressler et al.\\ 1997).  Detailed photometric observations of other intermediate redshift ($z \\simless 0.4$) clusters have confirmed the original results of Butcher \\& Oemler (e.g.\\ Millington \\& Peach 1990; Luppino et al.\\ 1991; Rakos \\& Schombert 1995).  Even though these clusters show an increased fraction of blue galaxies, they still contain a population of E/S0s distinguished by extremely red colors and a tight color--magnitude (CM) relation (a ``red envelope''). Both the mean color and the CM relation are consistent with that of present--day ellipticals (e.g.\\ Sandage 1972; Butcher \\& Oemler 1984; Arag\\'on-Salamanca et al.\\ 1991; Luppino et al.\\ 1991; Molinari et al.\\ 1994; Dressler et al.\\ 1994; Smail, Ellis \\& Fitchett 1994; Stanford, Eisenhardt \\& Dickinson 1995).  At redshifts of $z \\simgreat 0.4$, the red envelope has moved bluewards with redshift (Arag\\'on-Salamanca et al.\\ 1993; Smail et al.\\ 1994; Rakos \\& Schombert 1995; Oke, Gunn \\& Hoessel 1996; Lubin 1996; Ellis et al.\\ 1997; Stanford, Eisenhardt \\& Dickinson 1997). At $z \\sim 0.9$, there are few cluster members with colors nearly as red as present--day ellipticals. The color distribution of this high-redshift elliptical population is relatively narrow, and the trend is uniform from cluster to cluster; this suggests a homogeneous population which formed within a narrow time span (e.g.\\ Bower, Lucey \\& Ellis 1992a,b). Dickinson (1995) finds similar results in a cluster of galaxies which is associated with the $z = 1.206$ radio galaxy 3C 324. The galaxies in this cluster exhibit a narrow, red locus in the CM magnitude diagram. This branch is $\\sim 0.6$ mag bluer than the expected ``no--evolution'' value, though the intrinsic rms color scatter is only 0.2 mag. The observed color trend for the red envelope of galaxies in this data is consistent with passive evolution of an old stellar population formed by a single burst of star formation at redshifts of $z \\simgreat 2$.  The reasonably small color scatter would imply closely synchronized intra--cluster star formation (Bower et al.\\ 1992a,b; Arag\\'on-Salamanca et al.\\ 1993; Dickinson 1995; Ellis et al.\\ 1997; Stanford et al.\\ 1997).  The high-resolution imaging of HST has been essential in understanding the evolutionary processes occurring at intermediate redshifts (see e.g.\\ Abraham et al.\\ 1996a). Morphological classifications can be made on scales of $\\sim 1~{\\rm kpc}$, providing a direct comparison with ground-based classifications of nearby galaxies. A comprehensive survey of 10 clusters of galaxies at $z = 0.37 - 0.56$ has revealed a significant change relative to local clusters in the composition and behavior of the galaxy populations (Smail et al.\\ 1997, hereafter S97; Dressler et al.\\ 1997, hereafter D97). The authors have visually classified over 6000 galaxies based on the Revised Hubble Scheme used to classify nearby galaxies (e.g.\\ Sandage 1961; Sandage \\& Bedke 1994). These classifications are used to quantify the morphological composition of each cluster. Their results indicate that the morphology--density relation is qualitatively similar to that in the local universe in those intermediate redshift clusters which are centrally-concentrated and compact; however, the relation is non-existent in the loose, open clusters. Even so, all of the clusters exhibit a roughly similar make-up of galaxy morphologies.  The fraction of ellipticals is the same or larger than that in local clusters; the S0 fraction, however, is $\\sim 2-3$ times lower, with a corresponding increase in the cluster spiral population. These findings imply that the elliptical population is already in place by $z \\sim 0.5$, but a large fraction of the S0 galaxies are formed between redshifts of $z \\sim 0.5$ and $z = 0$ (D97). However, it should be noted that these classifications are typically derived from images which are not of comparable quality to the local data. Because of such uncertainty, the observed evolution in the S0 population is still in contention (e.g.\\ Stanford et al.\\ 1997; Andreon, Davoust \\& Helm 1997; Andreon 1998).  Because there appears to be significant evolution occurring between redshifts of $z \\sim 0.5$ and the present epoch, it is critical to extend these detailed observations to even higher redshifts if we are to understand the formation of galaxy morphology, as well as the mechanisms and timescales of this evolution. Therefore, we have undertaken an extensive observational program to study nine candidate clusters of galaxies at redshifts of $z > 0.6$. The cluster sample was chosen from the Gunn, Hoessel \\& Oke (1986) survey and the Palomar Distant Cluster Survey (PDCS; Postman et al.\\ 1996). For each cluster, we are in the process of obtaining deep $BVRI$ photometry from Keck and deep $K$ photometry from the KPNO 4-meter, low-resolution spectra from Keck, and high angular resolution imagery from HST. The observations and data processing procedures of this survey are the subject of the first paper in this series (Oke, Postman \\& Lubin 1998; hereafter Paper I).  The first two clusters to be completed in this observational program are CL0023+0423 and CL1604+4304 at redshifts of $z = 0.84$ and $z = 0.90$, respectively (see Paper I). In this paper, we have used HST images to undertake a detailed morphological analysis of the galaxy populations in the central regions of these two clusters. The reduction and analysis of the Keck $BVRI$ photometry and spectra of the galaxies in these cluster fields are discussed in the second paper of this series (Postman, Lubin \\& Oke 1998; hereafter Paper II). However, the galaxy parameters presented in Paper II are used in this paper, specifically for a comparison with the morphological properties.  In Sect.\\ 2, we provide a brief description of the data. In Sects.\\ 3 and 4, we describe the automated and visual galaxy classification procedures used in this paper and present a comparison between the two techniques.  In Sect.\\ 5, we examine the morphologies of the galaxies in the two cluster fields. This includes the relationship between morphology and other galaxy properties, as well as the overall distribution of morphologies in the cluster. A discussion and summary of our conclusions are presented in Sects.\\ 6 and 7. In the following analyses, we have assumed $q_{0} = 0.1$ (e.g\\ Carlberg et al.\\ 1996) and $H_{0} = 100~h~{\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "As part of an observational program to study distant clusters of galaxies, we have examined the morphological properties of the galaxies in two cluster fields, CL0023+0423 at $z = 0.84$ and CL1604+4304 at $z = 0.90$, using high-resolution HST images. The morphology of the individual galaxies have been studied by two methods; 1) a quantitative description of the structural properties of $\\sim 600$ galaxies per cluster field is provided by the Medium Deep Survey automated data reduction and ``bulge+disk'' object classification software; 2) the brightest subsample of $\\sim 200$ galaxies per cluster field are assigned a more detailed morphological description through a visual classification based on the revised Hubble scheme. A comparison between the two techniques shows that there is a reasonable correlation between the parameters of the automated and visual classifications (see also Lahav et al.\\ 1995). To investigate the morphological composition of the two galaxy clusters, we have used the visual classifications of the brightest subsample of galaxies in each field.  Our main conclusions are summarized below.  \\newcounter{discnt}  \\begin{list} {\\arabic{discnt}.}  {\\usecounter{discnt}}  \\item The color-magnitude diagrams and the color histograms of all (field + cluster) galaxies in the two cluster field show a clear progression in color between early- and late-type galaxies. As expected, the elliptical and S0 galaxies are redder, on average, than the spirals and irregulars. This trend is also apparent in the color ages which represent the time since the last period of major star formation. 83\\% of the galaxies classified as late-type (spiral or irregular/peculiar) have color ages less than 2 Gyr. In contrast, 55\\% of the galaxies classified as early-type have color ages of greater than 2 Gyr, and 73\\% of all galaxies with ages greater than 3 Gyr are classified as early-type. In addition, cluster galaxies are typically older than field galaxies at similar redshifts. This is due in large part to the fact that there are more early-type galaxies in these systems.  \\item We see a distinct correlation between the galaxy morphology and the corresponding spectral features. We have specifically examined this relation for those galaxies which are confirmed cluster members. The majority of galaxies that are visually classified as ellipticals show spectra which are typical of nearby, red elliptical galaxies. However, some of the galaxies visually classified as ellipticals turn out to be blue compact galaxies with spectra characterized by fairly strong, narrow emission lines. Normal late-type galaxies typically have spectra with blue colors and [\\ion{O}{2}] emission, while the presence of strong star-formation features, such as extremely high equivalent width [\\ion{O}{2}], ${\\rm H\\beta}$, and/or [\\ion{O}{3}] emission, is always accompanied by peculiar morphologies which suggest recent mergers or interactions.  \\item Despite being at very similar redshifts, the two cluster systems contain very different galaxy populations as indicated by their background-subtracted morphological distributions.  We have examined all galaxies brighter than $M_{V} = -19.0 + 5~{\\rm log}~h$ in the central $\\sim 0.5~h^{-1}~{\\rm Mpc}$ of the cluster.  CL0023+0423 has a galaxy population which is more typical of groups and the field population. The numbers from the statistical distribution are consistent with almost all of the galaxies being normal spirals. The spectral results support these numbers, independently indicating spiral fractions of 66\\% or more.  CL1604+4304, in contrast, has a morphological composition which is characteristic of a normal, present-day rich cluster. Early-type galaxies comprise 76\\% of all galaxies in this region. In this population, the ratio of S0 galaxies to ellipticals is $1.7 \\pm 0.9$, consistent with galaxy populations found in local clusters (Dressler 1980a).  \\item The morphological results support the conclusions of the dynamical analysis presented in Paper II. CL0023+0423 is apparently two independent systems separated in velocity by $\\sim 2900~{\\rm km~s^{-1}}$.  The velocity dispersions and implied masses indicate that these systems are similar to local galaxy groups or poor clusters. Though this may be a chance projection, the dynamical and morphological evidence may indicate that we are seeing the merger of two spiral-dominated galaxy groups (see Lubin, Postman \\& Oke 1998a). The velocity histogram of CL1604+4304, on the other hand, is consistent with a Gaussian distribution, implying that this system formed at an earlier epoch and is already relaxed. The velocity dispersion and implied mass of this system are consistent with an Abell richness class 2 or 3 cluster.  \\end{list}   \\vskip 0.5cm  We thank the anonymous referee for his thorough review of this paper. Alan Dressler, Chris Fassnacht, and Ian Smail are thanked for useful discussions, comments, and material aids to this paper.  It is also a great pleasure to thank Allan Sandage for his generous gift of time and invaluable expertise to this project.  The W.M. Keck Observatory is operated as a scientific partnership between the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration.  It was made possible by generous financial support of the W. M. Keck Foundation. LML graciously acknowledges support from a Carnegie Fellowship.  This research was supported in part by {\\it HST} GO analysis funds provided through STScI grant GO-06000.01-94A and {\\it HST} Archival grant 7536.   \\clearpage"
    },
    "9804/astro-ph9804208_arXiv.txt": {
        "abstract": "s{Galactic synchrotron emission is a potentially confusing foreground, both in total power and in polarization, to the Cosmic Microwave Background Radiation. It also contains much physical information in its own right. This review examines the amplitude, angular power spectrum and frequency spectrum of the synchrotron emission as derived from the presently available de-striped maps. There are as yet no maps at arcminute resolution at frequencies above 2.4 GHz.  This incomplete information is supplemented with data from supernovae, which are thought to be the progenitors of the loops and spurs found in the Galactic emission. The possible variations of the frequency spectral index from pixel to pixel are highlighted.  The relative contributions of free-free and synchrotron radiation are compared, and it is concluded that the free-free contribution may be smaller than had been predicted by COBE.  New high resolution polarization surveys of the Galactic plane suggest detail on all scales so far observed.  At high latitudes the large percentage polarisation means that the foreground contamination of the polarised CMB signal will be more serious than for the unpolarized radiation.} ",
        "introduction": "Galactic emission at radio wavelengths is important to understand in its own right.  Moreover it is crucial to be able to quantify and remove this component as a foreground to the cosmic microwave background (CMB).  Both synchrotron and free-free emission contribute to this foreground, with the synchrotron emission dominating at low frequencies ($\\leq$1 GHz).  The synchrotron emissivity is a function of both the relativistic (cosmic ray) density and the local magnetic field strength.  The luminosity at frequency $\\nu$ is given by  \\be I(\\nu) = const \\;  L N_0 B^{(p+1)/2} \\nu^{-(p-1)/2} \\ee  \\noindent where $N_0$ is the density of relativistic electrons, $L$ is the emission depth, $B$ is the magnetic field and the relativistic electron energy spectrum\\cite{ve} is given by $dN/dE = N_0 E^{-p}$.  The radio spectral index is $\\alpha= (p-1)/2$ in energy terms or $2+\\alpha$ when expressed as a brightness temperature $T_B \\propto \\nu^{-(2+\\alpha)} = \\nu^{-\\beta}$.  Within the interstellar magnetic field of 2 to 5 microgauss, emission at GHz frequencies is characteristically from relativistic electrons with an energy of 1 to 10 GeV.  Both $B$ and $N_0$, as well as $p$, will vary from point to point in the Galactic disk and nearby halo.  The cosmic ray electrons are thought to originate mainly in supernovae then diffuse outwards in the expanding remnant.  Structure will be formed in the remnant as it collides with the non-uniform ambient medium.  The magnetic field will be likewise amplified in compression regions and vary in strength and direction.  The net effect is to produce elongated synchrotron emission structures on a wide range of scales.  The spectral index of the emission will vary with position for two reasons.  Firstly the electron spectral index varies from one supernova to another and secondly the spectrum steepens ($\\Delta p = +1$) with time due to radiation energy loss thus giving an age-dependent spectral index.  This paper will describe the synchrotron features in and near the Galactic plane which are believed to give rise to the structures seen at higher galactic latitudes.  All-sky and large area surveys are assessed to give information about the amplitude and spectrum of the high latitude emission which is a potential confusing foreground to the CMB.  Comments are given about the role of synchrotron polarization and of free-free emission. ",
        "conclusions": ""
    },
    "9804/astro-ph9804102_arXiv.txt": {
        "abstract": "We report on {\\it ROSAT} HRI observations of the $z=0.61$ radio galaxy 3C\\,220.1. The X-ray emission from this object consists of an extended component, which we attribute to luminous cluster emission, and a compact central source. The compact component is too bright to be modelled as a cooling flow under some plausible assumptions for the hot gas temperature and distribution; we suggest instead that it is directly related to the core of the radio source. The X-ray flux of the compact component is consistent with the prediction of Worrall \\etal\\ (1994) that all powerful radio galaxies should have a central jet-related X-ray emission component that is proportional in strength to the radio core flux density. Other observations of distant 3CR radio sources are consistent with this model. ",
        "introduction": "In unified models of powerful radio sources (e.g.\\ Barthel 1989) core-dominated quasars, lobe-dominated quasars and radio galaxies are the same objects, with the apparent differences attributed to the effects of relativistic beaming and anisotropic obscuration on a population of sources oriented randomly with respect to the line of sight. In quasars we see the AGN directly, while radio galaxies have their jet axis at large angles to the line of sight so that a torus of gas and dust obscures optical continuum and broad-line emission from the AGN. This torus should also obscure soft X-ray emission originating close to the central engines, leading to suggestions (e.g.\\ Crawford \\& Fabian 1996) that X-ray emission from powerful radio galaxies should in general be dominated by thermal emission from hot cluster gas. However, an additional X-ray component may arise from the radio-emitting plasma directly, through synchrotron or synchrotron self-Compton radiation, or indirectly, through mechanisms such as the inverse-Compton scattering of external photons. If the different components can be separated, X-ray observations can provide physical insights into the active nucleus, jets and large-scale environment.  There is substantial evidence that unabsorbed radio-related non-thermal X-rays are seen. Fabbiano \\etal\\ (1984) first stressed that a correlation between the total soft X-ray and radio-core luminosity in radio galaxies implied a nuclear, jet-related origin for at least some of the X-ray emission. More recent work has strengthened this conclusion, as high-resolution X-ray observations have allowed point-like and extended components to be separated (Worrall \\& Birkinshaw 1994; Edge \\& R\\\"ottgering 1995; Worrall 1997).  Although component separation is generally better for closer, less powerful radio galaxies, compact soft X-ray emission is also seen in powerful narrow-line FRII radio galaxies that are known to lie in sparse environments, where cluster emission is not a source of confusion (e.g.\\ Hardcastle, Birkinshaw \\& Worrall 1998a). The full complexity is illustrated by X-ray observations of the nearby powerful cluster radio galaxy Cygnus A. {\\it EXOSAT} and {\\it Ginga} have found evidence for highly obscured core emission, ($N_H \\sim 4 \\times 10^{23}$ cm$^{-2}$; Arnaud \\etal\\ 1987, Ueno \\etal\\ 1994) which should not have been seen with {\\it ROSAT}, given its low-energy X-ray passband. However, an unresolved core component was detected with the {\\it ROSAT} HRI (Harris, Perley \\& Carilli 1994), implying that Cygnus A's core has both an absorbed and an unabsorbed X-ray component; while the absorbed component may be associated with emission from the AGN, the unabsorbed component may be radio-related (Worrall 1997). Birkinshaw \\& Worrall (1993) argued in the case of NGC 6251 that a plausible source for such compact X-ray emission is synchrotron self-Compton (SSC) radiation from the base of the radio jet, originating on scales larger than that of the torus and so avoiding absorbtion. Such emission would be suppressed, but not eliminated, by relativistic beaming effects in radio galaxies and [as discussed in Worrall \\etal\\ (1994) and references therein] almost certainly dominates the X-ray emission in core-dominated quasars.  It is therefore important in interpreting the X-ray emission from radio galaxies to have spectral or spatial information capable of distinguishing a non-thermal or compact component from thermal or extended emission. High-redshift powerful radio galaxies are important as counterparts to quasars in unified models, and Worrall \\etal\\ (1994) performed such a spatial separation for the $z = 1$ radio galaxy 3C\\,280. They showed that the unresolved component fell on an extension of the correlation between X-ray and radio core flux observed for core-dominated quasars, consistent with the model discussed above. In this paper we report X-ray observations of the radio galaxy 3C\\,220.1 with the {\\it ROSAT} HRI. We use the high resolution of the HRI to constrain the contributions from compact and extended emission.  3C\\,220.1 is an FRII (Fanaroff \\& Riley 1974) narrow emission-line radio galaxy at $z=0.61$ (Spinrad \\etal\\ 1985). With $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0$, used throughout the paper, its 178-MHz luminosity is $3.6 \\times 10^{27}$ W Hz$^{-1}$ sr$^{-1}$. Radio imaging (Burns \\etal\\ 1984; Jenkins, Pooley \\& Riley 1977; Harvanek \\& Hardcastle 1998) shows it to be a typical classical double object with largest angular size 35 arcsec (see Fig.\\ \\ref{contour}, inset); at this redshift one arcsecond corresponds to 8.94~kpc, so the projected linear size of 3C\\,220.1 is about 300 kpc. Burns \\etal\\ (1984) report an unusually prominent one-sided jet in the eastern lobe, and the radio core is also comparatively prominent, which may be an indication that the source is significantly affected by relativistic beaming, although no broad emission lines are reported by Spinrad \\etal\\ (1985). The prominent radio core, compared with the weak cores of the objects observed by Worrall \\etal\\ (1994), was the motivation for the present observations, since it allows us to probe the possible core X-ray and radio association. Optical observations show no evidence for a rich cluster near 3C\\,220.1, although the presence of a gravitational lens arc with $z=1.49$ in HST observations implies a deep potential well (M.\\ Dickinson, private communication, 1997). ",
        "conclusions": "We observed 3C\\,220.1 in order to test the hypothesis of Worrall \\etal\\ (1994) that all powerful radio sources exhibit a compact X-ray component related to their radio core. We find that there is strong evidence for a compact X-ray component in this object, that it is too bright to be attributed to a cooling flow under some simple assumptions, and that it follows the expected positive correlation between radio and X-ray core flux or luminosity. 3C\\,220.1 lies slightly above the line of slope unity plotted through the core-dominated quasars (Fig.\\ \\ref{cores}). Together with its unusually prominent radio core and one-sided radio jet, this may be an indication that we are viewing the source at an angle to the line of sight which is close to the radio galaxy-quasar boundary, so that its central X-ray emission includes a non-jet-related component originating close to the AGN.  Extended X-ray emission is also unequivocally detected around 3C\\,220.1, with core radius and luminosity comparable to that of nearby rich clusters. This is qualitatively consistent with the HST detection of a gravitational lensing arc near the source."
    },
    "9804/astro-ph9804334_arXiv.txt": {
        "abstract": "Deep infrared and optical images are presented of three proposed remnants of Thorne-$\\dot{\\rm Z}$ytkow Objects (T$\\dot{\\rm Z}$O). In particular, the IR data go several infrared magnitudes deeper than previous observations and in at least one case reveal the existence of weak objects within the error circles. It is argued, however, that none of the objects is likely to be the binary companion to the X-ray source in that region. These data present severe limits on any possible star or residual envelope at the distances of the respective X-ray objects. ",
        "introduction": "Recent work by several authors, especially van Paradijs et al. (1995) and Mereghetti \\& Stella (1995), have identified a group of pulsating X-ray sources that appear to have no detectable optical counterpart. In addition, these objects show no evidence of any binarity and all have pulse periods in the 5-10s range. Their X-ray luminosities are all sufficiently high that accretion must be occurring on to the neutron star from some source other than the interstellar medium. One possible scenario is that they are the first objects identifiable as the remains of Thorne-$\\dot{\\rm Z}$ytkow Objects (T$\\dot{\\rm Z}$Os).  The progenitor T$\\dot{\\rm Z}$Os would have consisted of a neutron star embedded in the core of another star (Thorne \\&$\\dot{\\rm Z}$ytkow 1977).  It is thought that these T$\\dot{\\rm Z}$Os could be formed as a result of the evolution of high mass systems through a common envelope phase (Ghosh et al. 1997). Once formed, thermonuclear burning via the {\\it r-p} process occurs in the core region of the T$\\dot{\\rm Z}$O. This supports the highly convective outer envelope against gravitational forces.  However, once there is insufficient fuel remaining for the {\\it r-p} process to occur efficiently, the outer envelope collapses on the Kelvin-Helmholtz timescale (Podsiadlowski, Cannon and Rees, 1995). Some of the envelope material then forms into a disk around the neutron star with a radius that is dependent on the amount of angular momentum in the system. It is from this disk of material that the neutron star may accrete enough matter to power the X-ray pulsations seen.  In order to explore this possibility deep infrared imaging observations have been carried out from UKIRT and the NASA IRTF to search for any IR emission from such a disk around three of the neutron stars. The observations reach as deep as J$\\sim$20 and in some cases reveal the presence of faint objects in the X-ray error circles. However, the characteristics of the IR colours, and other considerations, strongly suggest that these are not the TZO envelopes around the neutron stars. However, the new stringent limits on any IR companion in these systems leave these objects as enigmatic as ever.  The three systems investigated here are:  \\subsection{1E2259+586}  The X-ray pulsar 1E2259+586 lies in the centre of the supernova remnat G109.1-1.0. The most recent estimates on its age make it approximately 3000 years old (Parmar {\\it et al.} 1997), though values up to 10000 years have been quoted. G109.1-1.0 is at an estimated distance of about 4 kpc. Fahlman{\\it et al.} (1982) demonstrated that the X-ray column densities to both the pulsar and the SNR were similar and hence that these two objects were probably related.  However, the latest results from Parmar et al.(1997) indicate that this may not be the case. They find the X-ray column density to the pulsar to be greater than that to the supernova remnant. They suggest that this could be due to the pulsar being at a greater distance than the supernova remnant or to the presence of absorbing material local to the pulsar.  However, assuming the same distance to 1E2259+586 as the SNR implies a pulsar X-ray luminosity of $\\approx10^{28}$ Watts.  Results from {\\it EXOSAT}, {\\it GINGA} and {\\it BeppoSAX} (Hanson et al. 1988, Koyama et al. 1989 and Parmar et al. 1997 respectively) have shown that the 7s period of 1E2259+586 is slowly increasing, but the spindown is too slow to power the observed luminosity.  There is no evidence of binary motion, the upper limit on $a_{x} \\sin i$ from pulse timing measurements being as low as 30 light-ms (Mereghetti, Israel and Stella, 1998).  Early searches for the optical counterpart (Fahlman {\\it et al.} 1982; Margon \\& Anderson 1983) found several possible candidates, one of which (star D) was tentatively identified with 1E2259+586 by Middleditch, Pennypacker \\& Burns (1983) on the basis of IR pulsations.  However, more sensitive fast photometric observations by Davies {\\it et al.} (1989) failed to find any pulsations, and there would now appear to be little positive evidence associating star D with 1E2259+586.  A fainter, multicolour photometric study of all the possible counterparts by Davies \\& Coe (1991), together with revised astrometry, updated the list of possible counterparts.  Coe \\& Jones (1992) presented the first optical spectroscopy of all the candidates brighter than V=23 and re-analysed the Einstein X-ray images of this source taken in 1981 in order to check the position and size of the X-ray error circle. They found that the relocation and reduction in the size of the error circle stemming from the re-processing of these data considerably changes the perspective on possible counterparts to the X-ray pulsar. In addition, a careful study of the images reveals previously unreported features that are possibly associated with jet activity from the pulsar, a possibility originally suggested by Gregory and Fahlman (1980). However, Hurford \\& Fesen 1995 report subsequent {\\it Rosat} HRI observations which, while confirming the existence of such features, they suggest are just statistical fluctuations in the SNR emission.  \\subsection{4U0142+62}  This persistent source has appeared in most X-ray catalogues and was originally confused (due to poor angular resolution) with the Rosat source RX J0146.9+6121 (=LSI +61 235) - the two sources are only 24 arcminutes apart. Measurements by White et al. (1987) with {\\it EXOSAT} reported a 25m modulation from the region (later shown to be coming from LSI +61 235 (Mereghetti, Stella \\& De Nile 1993)) but using the unmodulated lower energy data they were able to identify the position of 4U0142+62. The detection of X-ray pulsations at 8.7s was reported by Israel et al. (1994) who concluded that they were most likely coming from the optically unidentified system 4U0142+62 rather than the new {\\it Rosat} source. Further measurements by White et al. (1996) using {\\it ASCA} confirmed this conclusion and also demonstrated evidence for an excess in the X-ray halo around the source. This could be due to material local to the pulsar along the line of site, possibly in the form of a molecular cloud or perhaps the remains of a TZO envelope. No evidence of any binary motion has been reported from the X-ray signal.  The accurate position reported by White et al. (1987) allowed them to locate the X-ray error circles from both {\\it EXOSAT} and {\\it Einstein} on the sky. As in the case of 1E2259+586 there is no evidence of any obvious optical candidates within the defined regions. They obtained an R band image and set a limit of R$\\geq$22.5 for any counterpart. The ASCA observations (White et al. 1996) determined the column to the source to be 8 x 10${^{21}}$ cm${^{-2}}$ which in turn allows them to deduce an optical extinction of A$_{v}$$\\sim$4.7. The distance to the source is not well defined but probably lies in the range 0.5--2.0 kpc. As in the case of 1E2259+586 a molecular cloud lies near to, or in front of 4U0142+62 (Mereghetti \\& Stella 1995) and may be affecting the column measurement.  \\subsection{RX J1838.4--0301}  This source was originally reported by Schwentker (1994) as 5.5s pulsar possibly associated with a SNR at a distance of 4kpc. Interestingly, in this case the X-ray positional error circle was reported to include a $\\sim$14 mag optical candidate. Subsequently Mereghetti et al. (1997) reanalysed the same {\\it Rosat} data and reported are different interpretation of the data. They obtained an optical spectrum of the star in the error circle and concluded that this object was simply a K5 main sequence star and that the X-ray emission originated from its corona . ",
        "conclusions": "The first deep infrared and optical images have been presented of three proposed remnants of Thorne-$\\dot{\\rm Z}$ytkow Objects (T$\\dot{\\rm Z}$O). In particular, the IR data go several infrared magnitudes deeper than previous observations and in at least one case, 1E2259+586, reveal the existence of weak objects within the X-ray error circle. It is argued, however, that in that case none of the objects seen is likely to be the binary companion to the X-ray source in that region. These data present new limits on any possible IR counterpart to the respective X-ray objects."
    },
    "9804/astro-ph9804272_arXiv.txt": {
        "abstract": "Recently Alard\\markcite{alard1996} proposed to detect the shift of a star's image centroid, $\\delta x$, as a method to identify the lensed source among blended stars.  Goldberg \\& Wo\\'zniak\\markcite{goldberg1997} actually applied this method to the OGLE-1 database and found that 7 out of 15 events showed significant centroid shifts of $\\delta x \\gtrsim 0.2$ arcsec. The amount of centroid shift has been estimated theoretically by Goldberg.\\markcite{goldberg1997}  However, he treated the problem in general and did not apply it to a particular survey or field, and thus based his estimates on simple toy model luminosity functions (i.e., power laws). In this paper, we construct the expected distribution of $\\delta x$ for Galactic bulge events by using the precise stellar LF observed by Holtzman et al.\\markcite{holtzman1998} using HST. Their LF is complete up to $M_I\\sim 9.0$ ($M_V\\sim 12$), corresponding to faint M-type stars.  In our analysis we find that regular blending cannot produce a large fraction of events with measurable centroid shifts. By contrast, a significant fraction of events would have measurable centroid shifts if they are affected by amplification-bias blending. Therefore, Goldberg \\& Wo\\'zniak's measurements of large centroid shifts for a large fraction of microlensing events confirms the prediction of Han and Alard that a large fraction of Galactic bulge events are affected by amplification-bias blending. ",
        "introduction": "Experiments to detect Massive Astronomical Compact Halo Objects (MACHOs) by monitoring the light variations of stars undergoing gravitational microlensing events are being carried out by several groups (MACHO, Alcock et al.\\ 1997a; EROS, Ansari et al.\\ 1996; OGLE, Udalski et al.\\ 1997; DUO, Alard \\& Guibert \\ 1997).\\markcite{alcock1997a, ansari1996, udalski1997, alard1997b} Since the lensing probability for a single source star is very low, these searches are being conducted toward very dense star fields such as the Large Magellanic Cloud and Galactic bulge. While searches towards these crowded fields result in an increased event rate, it also implies that many of the observed light curves include light from unresolved stars that are not being lensed: the blending problem.   Depending on the source for the blended light, blending can be classified into several types. The first type, ``regular blending'', occurs when a bright source star registered on the template plate is lensed and its flux is blended with the light from numerous faint unresolved stars below the detection limit imposed by crowding. Regular blending affects the results of lensing experiments in various ways. First, it makes the measured event timescale shorter than the true one. Since the lens mass scales as $M \\propto t_{\\rm E}^{2}$, the lens mass determined from the measured timescales will be underestimated and the lens population will be misinterpreted. In addition, since the optical depth is directly proportional to the summation of event timescales, i.e., $\\tau \\propto \\sum_i t_{{\\rm E},i}$, the Galactic MACHO fraction determined from the optical depth without a proper blending correction is subject to great uncertainty (Di Stefano \\& Esin 1995).\\markcite{distefano1995}  The second type of blending occurs if the source for the blended light is the lens itself: ``lens blending'' (Kamionkowski 1995; Buchalter, Kamionkowski, \\& Rich 1996; Buchalter \\& Kamionkowski 1997; Alard 1997).\\markcite{kamionkowski1995, bucjalter1996, buchalter1997, alrad1997} Besides the effects of regular blending, lens blending has several additional effects on the result of lensing experiments. First, because detecting events due to lenses close to the observer is comparatively more difficult than detecting events produced by lenses near the source, lens blending makes the optical depth depend on the geometry of the lens system. As a result, the matter distribution derived from the optical depth distribution deviates from its true one. Secondly, lens blending causes the lensing optical depth to depend on the lens mass function since more massive stars, which contribute more to the total optical depth, tend to be brighter, resulting in a larger blending effect (Nemiroff 1997; Han 1998).\\markcite{nemiroff1997, han1998}   Finally, ``amplification bias'' blending occurs when one of several faint stars in the seeing disk below the detection limit is lensed, and the flux of the lensed star is associated with the flux from other stars in the integrated seeing disk (Nemiroff 1997).\\markcite{nemiroff997} In current experiments, photometry is carried out by comparing template images, in which only very bright stars are resolved and registered, with a series of images taken of the same field. The result is that in amplification-biased events the brightest star appears to be the source because the lensed star is too faint to be resolved. To be detected, the amplification-biased event must be highly amplified to overcome the high threshold flux from the brighter star. Therefore, the mean detection probability to detect events for each source star will be very low. However, if these faint stars comprise a significant fraction of the total number of stars, a considerable fraction of events might be amplification-biased (Han 1997b; Alard 1997a).\\markcite{han1997b, alard1997a} Moreover, due to the large amount of blended light from a bright star, the effects of blending for these events would be much more significant than those caused by other blending types.   There have been several methods proposed to correct for the blending problem. The first method is to introduce an additional lensing parameter representing the residual flux from unresolved faint stars into the light curve fitting process. However, this method suffers from large uncertainties in the derived lensing parameters as a result of parameter degeneracies (Wo\\'zniak \\& Paczy\\'nski 1997).\\markcite{wozniak1997} Early-warning systems (for MACHO: Pratt et al.\\ 1996; for OGLE: Udalski et al.\\ 1994; for PLANET: Albrow et al.\\ 1995)\\markcite{pratt1996, udalski1994, albrow1995} allow one to construct lensing light curves with high time resolution and small photometric errors, enabling one to detect small color shifts caused by blending (Buchalter et al.\\ 1996)\\markcite{buchalter996}. However, due to the narrow distribution of colors for Galactic bulge stars, the expected color shifts are small. Han (1997b)\\markcite{han1997} proposed to use the Hubble Space Telescope (HST) to provide blending corrections. By using the high resolving power of HST combined with color information from ground-based observations, one can identify the lensed source star in the blended seeing disk, thus the uncertainty in the derived timescale can be significantly reduced. However, this method requires costly HST times. One can also correct the blending effect statistically if the luminosity function(LF) of stars well below the detection limit can be constructed (Alcock et al.\\ 1997b)\\markcite{alcock1996}, but in this case we lose information about individual events.  Recently Alard (1996)\\markcite{alard1996} proposed to detect the shift of a star's image centroid, $\\delta x$, as a method to identify the lensed source among blended stars. Goldberg \\& Wo\\'zniak (1997)\\markcite{goldberg1997} actually applied this method to the OGLE-1 database and found that 7 out of 15 events showed significant centroid shifts of $\\delta x \\gtrsim 0.2$ arcsec. The amount of centroid shift has been estimated theoretically by Goldberg (1997).\\markcite{goldberg1997} However, he treated the problem in general and did not apply it to a particular survey or field, and thus based his estimates on simple toy model luminosity functions (i.e., power laws). In this paper, we construct the expected distribution of $\\delta x$ for Galactic bulge events by using the precise stellar LF observed by Holtzman et al.\\ (1998)\\markcite{holtzman1998} using HST. Their LF is complete up to $M_I\\sim 9.0$ ($M_V\\sim 12$), corresponding to faint M-type stars. In our analysis we find that regular blending cannot produce a large fraction of events with measurable centroid shifts. By contrast, a significant fraction of events would have measurable centroid shifts if they are affected by amplification-bias blending. Therefore, Goldberg \\& Wo\\'zniak's measurements of large centroid shifts for a large fraction of microlensing events confirms the prediction of Han (1997a) and Alard (1997) that a large fraction of Galactic bulge events are affected by amplification-bias blending. ",
        "conclusions": "The fact that events affected by amplification-bias blending produce large $\\delta x$ while the centroid shifts caused by regular blending are small can be understood analytically in the following way. To produce centroid shifts large enough to be measured, events should satisfy two conditions. First, the event should be highly amplified. For a very low amplification event ($A_{\\rm abs}\\sim 1$), the expected centroid shift will be small since $A_{\\rm abs}-1 \\sim 0$ and thus $\\eta \\sim 0$ in equation (1).  On the other hand, for a very high-amplification event ($A_{\\rm abs}\\sim \\infty$), one finds $\\delta x\\sim \\left\\vert\\langle {\\bf x}\\rangle -{\\bf x}_{j}\\right\\vert$ since the factor $\\eta$ approaches unity. However, not all events that satisfy the first condition produce large centroid shifts. The second condition is that the lensed star should be one of the faint stars in the blended seeing disk. If the lensed star is the brightest one in the effective seeing disk and its flux dominates the flux over those from other faint blended stars, i.e., $\\kappa\\sim 1$, the position of the CL before gravitational amplification will be very close to that of the lensed star, resulting in a small amount of shift ($\\left\\vert\\langle {\\bf x}\\rangle-{\\bf x}_j\\right\\vert \\sim 0$) since $\\sum_{i\\neq j} F_{0,i}{\\bf x}_i + F_{0,j}{\\bf x}_j\\sim F_{0,j}{\\bf x}_j$ and $\\sum_i F_{0,i} \\sim F_{0,j}$. On the other hand, if the lensed source is very faint, i.e., $\\kappa\\sim 0$, the light from the source star has negligible effect on the position of CL, resulting in high possibility of a large centroid shift.  In summary, the expected centroid shifts for various extreme cases of amplification and source star brightness are: $$ \\cases{ \\eta\\sim 0\\ {\\rm and}\\ \\delta x\\sim 0   & for a low amplification event \\cr & ($A_{\\rm abs}\\sim 1$) \\cr \\delta x\\sim \\left\\vert\\langle{\\bf x}\\rangle-{\\bf x}_j\\right\\vert\\sim 0 & for a high-amplification event\\cr & with luminous source \\cr & ($A_{\\rm abs}\\sim \\infty$, $\\kappa\\sim 1$) \\cr \\delta x\\sim \\left\\vert\\langle{\\bf x}\\rangle-{\\bf x}_j\\right\\vert & for a high-amplification event\\cr & with faint source \\cr & ($A_{\\rm abs}\\sim \\infty$, $\\kappa\\sim 0$). \\cr } \\eqno(4) $$  For amplification-biased events, source stars are in general very faint ($\\kappa\\sim 0$), mostly far below the detection limit. Despite their low luminosities, the fact that they are detected implies that the source stars are highly amplified ($A_{\\rm abs}\\sim \\infty$). Therefore, the conditions for large centroid shifts agree well with those for amplification-biased events. On the other hand, regular blended events do not meet these conditions. First, due to relatively small amount of blended light, regular blended events do not need to be highly amplified for detection ($A_{\\rm abs}\\sim 1$). Although they can be highly amplified ($A_{\\rm abs}\\sim \\infty$), the dominance of their fluxes ($\\kappa\\sim 1$) over those from other faint blended stars will result in small $\\delta x$.   As demonstrated by the large centroid shifts for a significant fraction of events, the effect of amplification-bias blending on the results of lensing experiments is important. However, the methods to correct for the blending effect mentioned in \\S\\ 1 have various limitations in application. One very simple but very practical method to minimize the effect of amplification-bias blending is to monitor only very bright stars. With increasing reference image brightness, the required amplification for detection becomes higher, resulting in a lower probability of amplification-biased events. To show how the blending effect decreases with increasing reference image brightness, we simulate Galactic bulge events which are expected to be detected for various threshold reference image brightnesses. For each event, we compute the light fraction of the source star $\\kappa$ and the timescale decrease factor $\\eta$.  The distributions of $\\kappa$ and $\\eta$ are shown in the upper panel of Figure 6. For a given fraction of source star flux, the observed timescale is reduced by $\\eta=t_{\\rm eff}/t_{\\rm E}= \\left[ 2(1-A_{\\rm min}^{-2})^{-1/2}-2 \\right]^{1/2}; \\ A_{\\rm min} = 0.34/\\kappa + 1$. Here events are assumed to be detected as long as they can amplify the reference image flux by more than a factor of 1.34. However, highly blended events will have short $t_{\\rm eff}$, resulting in low detectability. Therefore, we correct the distributions by the detection efficiency. We assume that the detection efficiency is linearly proportional to the timescale decrease factor, i.e., $\\epsilon \\propto \\eta$. The efficiency-corrected distributions $f(\\kappa)$ and $f(\\eta)$ are presented in the middle panels. In the lower panels we present the distributions of the fraction of events with $\\kappa \\ge \\kappa_{\\rm lim}$, $1-\\int_{0}^{\\kappa_{\\rm lim}} f(\\kappa)d\\kappa$, and $\\eta \\ge \\eta_{\\rm lim}$, $1-\\int_{0}^{\\eta} f(\\eta')d\\eta'$. From these distributions one finds that under the current threshold reference image brightness of $M_I\\sim 3$, the fraction of events with little blending effects ($\\kappa \\gtrsim 0.9$ or $\\eta \\gtrsim 0.9$) is $\\lesssim 10\\%$. However, as the brightness of the threshold reference image increases, this fraction gradually increases until it becomes $\\sim 80\\%$ when only stars brighter than $M_I\\sim 0$ are monitored, which corresponds to the brightness of Galactic bulge clump giant stars (Paczy\\'nski \\& Stanek 1998).\\markcite{paczynski1998} The MACHO group (1997a)\\markcite{alcock1997a} already applied this method and their optical depth determination is based on clump giant stars. However, by monitoring significantly fewer stars at a decreased event rate, the statistical precision of the lensing experiments will be lowered. A more general solution for the blending correction is provided by the rapidly progressing image subtraction technique (Alard \\& Lupton 1997b; Tomaney 1998)\\markcite{alard1997c, tomaney1998} which is also being applied to detect microlensing events towards M31 by the Colombia-Vatt group (Crotts \\& Tomaney 1996; Tomaney \\& Crotts 1996).\\markcite{crotts1996, tomaney1996}"
    },
    "9804/astro-ph9804050_arXiv.txt": {
        "abstract": "We investigate if the gamma ray halo, for which recent evidence has been found in EGRET data, can be explained by neutralino annihilations in a clumpy halo. We find that the measured excess gamma ray flux can be explained through a moderate amount of clumping in the halo. Moreover, the required amount of clumping implies also a measureable excess of antiprotons at low energies, for which there is support from recent measurements by the BESS collaboration.  The predicted antiproton fluxes resulting from neutralino annihilations in a clumpy halo are high enough to give an excess over cosmic-ray produced antiprotons also at moderately high energies (above a few GeV). This prediction, as well as that of one or two sharp gamma lines coming from annihilations into $\\gamma\\gamma$ or $Z\\gamma$ can be tested in upcoming space-borne experiments like AMS and GLAST. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804099_arXiv.txt": {
        "abstract": "We  investigate  the nature of stellar populations of major galaxy mergers between late-type spirals considerably abundant in interstellar medium by performing numerical simulations designed to solve both the dynamical and chemical evolution in a self-consistent manner. We particularly consider that the star formation history of galaxy mergers is a crucial determinant  for the nature of stellar populations of merger remnants, and therefore investigate how the difference in  star formation history between galaxy mergers affects the chemical evolution of galaxy mergers. We found that the rapidity of star formation, which is defined as the ratio of the dynamical time-scale to the time-scale of gas consumption by star formation, is the most important determinant for a number of fundamental characteristics of stellar populations of  merger remnants. Main results obtained in this study are the following five. (1) A galaxy merger with more rapid star formation  becomes elliptical with larger mean metallicity. This is  primarily because in the merger with more rapid star formation, a smaller amount of metal-enriched gas is tidally stripped away during merging and consequently a larger amount of the gas can be converted to stellar component. This result demonstrates that the origin of the color-magnitude relation of elliptical galaxies can be  closely associated with the details of merging dynamics which depends on the rapidity of star formation in  galaxy mergers. (2) Negative metallicity gradient fitted reasonably well by power-low can be reproduced by  dissipative galaxy mergers with star formation. The magnitude of metallicity gradient is larger for an elliptical galaxy formed by galaxy merging with less rapid star formation. (3) Absolute magnitude of metallicity gradient  correlates  with that of age gradient in  galaxy mergers in the sence that a merger remnant with steeper negative metallicity gradient is more likely to show steeper age gradient. (4) The outer part of stellar populations is both older and less  metal-enriched than nuclei in an elliptical galaxy formed by galaxy merging with less rapid star formation. Moreover, the metallicity of the  outer part of gaseous component for some models with less rapid star formation is appreciably smaller than that of stellar one. This result implies that the origin of metal-poor hot gaseous $X$-ray halo in real elliptical galaxies can be essentially  ascribed to the dynamics of dissipative galaxy merging. (5) Irrespectively of the rapidity of star formation, the epoch of galaxy merging affects both the mean stellar metallicity and mean stellar age of merger remnants: Later galaxy mergers are more likely to become ellipticals with both younger and more metal-enriched stellar populations. This result reflects the fact that in the later mergers, a larger  amount of more metal-enriched interstellar gas is preferentially converted into younger stars in the later star formation  triggered by galaxy merging. These five results clearly demonstrate that even the chemical evolution  of elliptical galaxies can be strongly affected by  the details of dynamical evolution of galaxy merging, which is furthermore determined  by the rapidity of star formation of galaxy mergers. In particular, tidal stripping of interstellar gas and total amount of gaseous dissipation during galaxy merging are  demonstrated to play a vital role in determining a number of chemical properties of merger remnants. Based upon these results, we adopt a specific assumption of the luminosity dependence of the rapidity of star formation and thereby  discuss how successfully the present merger model can reproduce a number of fundamental chemical, photometric, and spectroscopic characteristics  of elliptical galaxies. ",
        "introduction": "Elliptical galaxies have been generally considered to be old, coeval and homogeneous systems passively evolving after the single initial burst of star formation associated with dissipative galaxy formation. This classical picture of elliptical galaxy formation appears to have been supported by the considerably tight color-magnitude relation of elliptical galaxies ( Bower, Lucey, \\& Ellis 1992; Ellis et al. 1997) and by relatively smaller  redshift evolution  of photometric properties of elliptical galaxies (Arag$\\rm \\acute{o}$n-Salamanca et al. 1993; Franx \\& van Dokkum 1996). A growing number of recent observational results, however, shed a strong doubt on this long-standing view of elliptical galaxy formation, and suggest that there is  great variety of star formation history between elliptical galaxies, such as the epoch of major star formation, the duration and efficiency of star formation (Worthey, Faber, \\& Gonzalez 1992; Matteuchi 1994; Faber et al. 1995; Bender 1996; Worthey, Trager, \\& Faber 1996). This tendency that elliptical galaxies show diversity in star formation history and nevertheless can actually keep the tightness of the color-magnitude relation is considered to be quite mysterious and thus to provide any theoretical models with a valuable insight on the elliptical galaxy formation. Such kind of mysterious nature observed in elliptical galaxies is demonstrated to hold equally for the dynamical  and kinematical properties of elliptical galaxies. For example, considerably small thickness of the fundamental plane of elliptical galaxies implies a rather smaller range of admitted dynamical state of the galaxies (Djorgovski \\& Davis 1987; Dressler et al. 1987 ; Djorgovski, Pahre, \\& de Carvalho 1996) whereas the morphological dichotomy between boxy-disky elliptical galaxies (Kormendy \\& Bender 1996) and the projected density profile systematically departing  from de Vaucouleurs $R^{1/4}$ law (Caon, Capaccioli, \\& D'Onofrio 1993) show  a great variety of major orbit families consisting the galaxies. These fundamental characteristics that elliptical galaxies show both diversity and uniformity in their chemical, photometric and dynamical properties have imposed  some stringent  but valuable constraints on any theoretical models of elliptical galaxy formation. What is the most vital in challenging the origin of elliptical galaxy formation in this kind of situation is to investigate whether or not both the chemical and photometric properties and dynamical and kinematical ones can be reproduced successfully by a specific model of galaxy formation in a reasonably self-consistent manner. The previous theoretical models addressing this important issue on elliptical galaxy formation are divided basically into two categories: The dissipative galactic collapse model (e.g., Larson 1976; Carlberg 1984) and the galaxy merger model (e.g., Toomre \\& Toomre 1972). As is suggested by Kormendy \\& Sanders (1992), these two dominant and apparently competing scenarios for elliptical galaxy formation are now converging, thus it would be crucial  to construct  one more realistic and sophisticated model of elliptical galaxy formation. Although there are a large number of important studies exploring the origin of elliptical galaxy formation along the dissipative collapse scenario, especially in the context of the nature of stellar populations (e.g., Arimoto \\& Yoshii 1987), we here restrict ourselves to the merger scenario  of elliptical galaxy formation.   Recent extensive studies of merger models of elliptical galaxy formation, mostly based upon numerical simulations, $appear$ to have succeeded in resolving most of the outstanding problems related to dynamical and kinematical properties of elliptical galaxies, such as the phase space density (Ostriker 1980; Carlberg 1986) and kinematical misalignment (Franx, Illingworth, \\& de Zeeuw 1991; Barnes 1992), by invoking the inclusion of bulge component, gaseous dissipation, and multiplicity of galaxy merging (Hernquist, Spergel, \\& Heyl 1993; Weil \\& Hernquist 1996; Barnes \\& Hernquist 1996). Although it would be safe to say that the galaxy merging between two late-type spirals  is one of the most promising candidates explaining more clearly the origin of elliptical galaxies at least in the context of the $dynamical$ $and$ $kinematical$ $properties$, however, there still remain  a number of unresolved  and apparently serious problems concerning the merger model (e.g., van den Bergh 1995). One of the most crucial problems among these is on whether  the fundamental characteristics of stellar populations of elliptical galaxies can be reproduced reasonably well by galaxy merging between two late-types spirals. Surprisingly, there are only a few works addressing this critical issue for the merger model, probably because it is considered to be rather difficult to solve the chemical evolution of galaxy mergers in which a number of competing physical processes are expected to affect strongly the chemical evolution of galaxy mergers. White (1980) and Mihos \\& Hernquist (1994) found  that the stellar populations of progenitor disks are not  mixed so well even by the violent relaxation during galaxy merging and consequently the metallicity gradient of progenitor disks is not so drastically washed out. The metallicity gradient of merger remnant is furthermore found not to  be fitted by power law  observed in elliptical galaxies (Mihos \\& Hernquist 1994). Schweizer \\& Seitzer (1992) discussed whether or not the bluer integrated $UBV$ color of elliptical galaxies with morphologically fine structure can be explained by secondary starburst induced by  major disk-disk galaxy mergers. Kauffmann \\& Charlot (1997) construct  a semi-analytic model of elliptical galaxy formation, which is based upon the hierarchical clustering in CDM universe and includes  rather simple chemical enrichment  process, and thereby demonstrate that the origin of the color-magnitude relation of elliptical galaxies can be reproduced successfully even in the CDM model of galaxy formation (See also Baugh, Cole \\& Frenk 1996.). Thus, since  there are only a few works addressing chemical and photometric properties  for the merger model, it is essential for the merger model to investigate more throughly the fundamental chemical and photometric properties of merger remnants, including the origin of color-magnitude relation (Faber 1973; Visvanathan \\& Sandage 1977), age and metallicity gradient (Peletier et al. 1990; Davies et al 1991), $\\rm Mg_{2} - \\sigma$ relation (Burstein et al. 1988), age-metal-conspiracy in stellar populations (Faber et al. 1995; Worthey et al. 1996), luminosity dependence of the line ratio [Mg/Fe] (Worthey et al. 1992),  metal-poor gaseous $X$-ray halo (Matsumoto et al. 1997), and the substantially metal-enriched galactic nuclei at higher redshift (Hamann \\& Ferland 1993).  What should be recognized foremost in investigating the nature of stellar populations in  merger remnants is that a glowing number of observational results have been accumulated which suggest the relatively earlier formation of elliptical galaxies. Tightness of the color-magnitude relation in the  cluster of galaxies (Bower et al. 1992, Ellis et al 1996), relatively smaller photometric evolution of cluster ellipticals (Arag$\\rm \\acute{o}$n-Salamanca et al. 1993), and the redshift evolution of the fundamental plane (Franx \\& van Dokkum 1996) all suggest the $typical$ formation epoch of elliptical galaxies is  earlier  than 2 in redshift. Furthermore, as is suggested by Kormendy \\& Sanders (1992), the fact that no galaxy in the $K$-band survey of Cowie et al. (1994) shows the global color resembling that of the Arp 220, which is considered to be  ongoing mergers and forming ellipticals, implies that the formation epoch of elliptical galaxies should be earlier than 1.0 in redshift. Silva \\& Bothun (1997) revealed that the fraction of mass of stellar populations with intermediate age to total mass in elliptical galaxies with morphologically fine structure is less than 15 percent. These results imply that if  elliptical galaxies are formed by galaxy merging, the epoch of galaxy merging should be relatively earlier and furthermore that the precursor disks of galaxy mergers may be extremely abundant in interstellar medium compared with the present spirals. Recent high quality imaging using $Hubble$ $Space$ $Telescope$ ($HST$)  has revealed that a  larger number of galaxies at faint magnitude are interacting/merging galaxies (e.g., van den Bergh et al. 1996), indicating furthermore that the potential candidate for elliptical galaxies formed by galaxy merging  are ubiquitous in higher redshift universe. Hence it is quite  reasonable and essential to study the nature of stellar populations of higher redshift  galaxy mergers between  disk galaxies with the gas mass fraction larger than 0.2, which is a typical value of the present late-type spirals, and thereby  to confirm whether or not elliptical galaxies can be  formed $actually$  by galaxy merging.   The purpose of this paper is to explore the nature of the stellar populations of a gas-rich disk merger which is considered to  be occurred the most frequently in the high redshift universe. We particularly investigate how successfully galaxy mergers between gas-rich spirals can reproduce a number of fundamental chemical, photometric, and spectroscopic  properties of elliptical galaxies. The layout of this paper is as follows. In \\S 2, we summarize  numerical models used in the present study and describe in detail   methods for analyzing the stellar populations produced by dissipative galaxy mergers with star formation. In \\S 3, we demonstrate how  a number of fundamental  characteristics of stellar populations in merger remnants are affected by  the star formation history of dissipative galaxy merging. In \\S 4, we discuss how successfully the present merger model can reproduce a number of observational results concerning the chemical, photometric, and spectroscopic properties of elliptical galaxies. In this section, we also point out the advantages and disadvantages of galaxy mergers in explaining both the chemical, photometric, and spectroscopic properties and dynamical and kinematical ones in real elliptical galaxies. The conclusions of the present study are given in \\S 5. ",
        "conclusions": "Main results obtained in the present  study are summarized as follows.   (1) Galaxy mergers with more rapid star formation  become ellipticals with larger mean stellar metallicity, primarily because in the mergers with more rapid gas consumption, a smaller  amount of metal-enriched gas is tidally stripped away during merging and consequently a larger  amount of the gas can be converted into stellar component. This result is demonstrated not to depend so strongly on the  other parameters such as the orbit configuration of galaxy merging and multiplicity of the mergers. These results suggest that the origin of the color-magnitude relation of elliptical galaxies can be  closely associated with the details of merging dynamics which depends on the rapidity of star formation (thus on the galactic luminosity) in  galaxy mergers.  (2) Negative metallicity gradient fitted reasonably well by power-low can be reproduced by  dissipative galaxy mergers with star formation, which is in good  agreement with the recent observational  results of elliptical galaxies. The absolute magnitude of metallicity gradient in each merger remnant depends  on the orbit configuration of each galaxy merging, suggesting that the observed dispersion in the absolute magnitude of metallicity gradient for a given luminosity range of elliptical galaxies reflects the diversity in the orbit configuration  of galaxy merging.  (3) Absolute magnitude of metallicity gradient  correlates  with that of age gradient in a merger in the sence that a merger remnant with steeper negative metallicity gradient is more likely to show steeper age gradient. This result reflects the fact  that the degree of violent relaxation and gaseous dissipation during merging strongly affect both the  age gradient and metallicity one.  (4) The outer part of stellar populations is both older and less  metal-enriched than nuclei in an elliptical galaxy formed by galaxy merging with less rapid star formation. Moreover galaxy mergers with less rapid star formation are more likely to become ellipticals with metal-poor gaseous halo. This result suggests that the formation of metal-poor $X$-ray halo actually observed in elliptical galaxies can be essentially ascribed to the dissipative galaxy merging between late-type spirals, and furthermore provides a clue to a solution for the iron abundance discrepancy problem in elliptical  galaxies.    (5) The epoch of galaxy merging affects both the mean stellar metallicity and the mean stellar age in merger remnants: Later galaxy mergers become ellipticals with both younger and more metal-enriched stellar populations. This result suggests that the origin of Worthey's 3/2 rule (Worthey et al. 1996), which is invoked in maintaining the tightness of the color-magnitude relation of elliptical galaxies, can be understood in terms of the difference in the epoch of galaxy formation and transformation, that is, the epoch of galaxy merging, between elliptical galaxies.   (6) Luminosity dependence of chemical, photometric, and spectroscopic properties in merger remnants, which is derived by adopting a specific assumption on the luminosity dependence of the rapidity of star  formation of galaxy mergers, does not match so reasonable well with that observed in real elliptical galaxies. This result implies that other fundamental physical processes expected to be dependent on the galactic luminosity should be incorporated into the present merger model for more successful comparison with observational trends of luminosity-dependent chemical, photometric, and spectroscopic properties of elliptical galaxies.   (7) As is described in the above (1) - (6), the details of gas dynamics of galaxy merging, in particular, the tidal stripping of metal-enriched interstellar gas and the degree of gaseous dissipation during merging, both of which depend on the star formation history of galaxy mergers, are demonstrated to determine even the chemical and photometric properties of merger remnants. These results can not be obtained until both the chemical and dynamical evolution during galaxy merging are solved numerically in a reasonably self-consistent way."
    },
    "9804/astro-ph9804320_arXiv.txt": {
        "abstract": "Physical properties of the atomic gas in spiral galaxies are briefly considered. Although both Warm (WNM, 10$^4$~K) and Cool (CNM, $\\sim$~100 K) atomic phases coexist in many environments, the dominant mass contribution within a galaxy's star-forming disk (R$_{25}$) is that of the CNM. Mass fractions of 60 to 90\\% are found within R$_{25}$. The CNM is concentrated within moderately opaque filaments with a face-on surface covering factor of about 15\\%. The kinetic temperature of the CNM increases systematically with galactocentric radius, from some 50 to 200~K, as expected for a radially declining thermal pressure in the galaxy mid-plane. Galaxies of different Hubble type form a nested distribution in T$_K$(R), apparently due to the mean differences in pressure which result from the different stellar and gas surface densities. The opaque CNM disappears abruptly near R$_{25}$, where the low thermal pressure can no longer support the condensed atomic phase. The CNM is apparently a prerequisite for star formation. Although difficult to prove, all indications are that at least the outer disk and possibly the inter-arm atomic gas are in the form of WNM, which accounts for about 50\\% of the global total. Median line profiles of the CNM display an extremely narrow line core (FWHM $\\sim$~6~km~s$^{-1}$) together with broad Lorentzian wings (FWHM $\\sim$~30~km~s$^{-1}$). The line core is consistent with only opacity broadening of a thermal profile. The spatial distribution of CNM linewidths is not random, but instead is extremely rich in structure. High linewidths occur in distinct shell-like structures with diameter of 100's of pc to kpc's, which show some correlation with diffuse H$\\alpha$ shells. The primary source of ``turbulent'' linewidth in the atomic ISM appears to be organized motions due to localized energy injection on a scale of a few 100 pc. ",
        "introduction": "Over the years a succession of eminent authors has considered the question of which physical processes determine the temperature of neutral atomic gas and which timescales are required to achieve (a local) thermodynamic equilibrium. This began with Field, Goldsmith and Habing (1969) and continued with Draine (1978), Shull and Woods (1985) and most recently Wolfire {\\it et al.} (1995). While some of the processes involved are quite straightforward, others, like the photoelectric emission from dust grains, have had to be substantially updated to reflect our growing knowledge of the properties and abundance of interstellar dust. The consensus which has emerged is that heating is in fact dominated by the photoelectric heating from small dust grains over a wide range of conditions and environments. Cooling, on the other hand, is regulated primarily by emission in the [CII] 158~$\\mu$m fine structure line at densities in excess of about 1~cm$^{-3}$ and by Ly$\\alpha$ emission at lower densities. The relative importance of the various mechanisms is nicely illustrated in Fig.~3 of Wolfire {\\it et al.} (1995). In the same figure can be seen the characteristic phase diagram for atomic gas. Two thermodynamically stable phases are found. The first is the so-called Warm Neutral Medium (WNM) which predominates at low densities and pressures, and has a kinetic temperature that rises toward lower pressures from a value of perhaps 5000~K to 10$^4$~K, with an accompanying increase in ionization fraction, $x$~$\\sim$~0.1 to 0.9. The second is the Cool Neutral Medium (CNM) which is primarily found at high densities and pressures, with a kinetic temperature that decreases towards higher pressures from a value of some 200~K to perhaps as low as 20~K. Over some range of intermediate pressures thought to be typical of the interstellar medium in the local neighbourhood of the Galaxy (P/k$\\sim$ 2000~K~cm$^{-3}$), the two phases can coexist in pressure equilibrium. However, given the strong gradient in the mid-plane thermal pressure which is likely to follow from the combination of a radial exponential stellar disk and a similarly declining gas surface density distribution, we can safely predict that the inner regions of disk galaxies will have predominantly rather cool CNM, while the outer regions of galaxies will eventually be dominated by the WNM. At intermediate radii we might also expect to see a radial increase in the CNM kinetic temperature in response to the declining mid-plane pressure. ",
        "conclusions": ""
    },
    "9804/astro-ph9804116_arXiv.txt": {
        "abstract": "Multidimensional cosmologies allow for variations of fundamental physical constants over the course of cosmological evolution, and different versions of the theories predict different time dependences.  In particular, such variations could manifest themselves as changes of the proton-to-electron mass ratio $\\mu=m_{\\rm p}/m_{\\rm e}$ over the period of $\\sim 10^{10}$ yr since the moment of formation of high-redshift QSO spectra.  Here we analyze a new, high-resolution spectrum of the $z = 2.81080$ molecular hydrogen absorption system toward the QSO PKS 0528--250 to derive a new observational constraint to the time-averaged variation rate of the proton-to-electron mass ratio. We find $| \\dot{\\mu}/ \\mu | < 1.5 \\times 10^{-14}$ yr$^{-1}$, which is much tighter than previously measured limits. ",
        "introduction": "The possibility of the variability of fundamental physical constants was first put forward by Dirac (1937) in the course of his discussion with Milne (1937).  Later it was considered by Teller (1948), Gamow (1967), Dyson (1972) and other physicists.  Interest in the problem increased greatly during the last decade, due to new developments in the Kaluza--Klein and supergravity models of unification of all the physical interactions. Chodos \\& Detweiler (1980), Freund (1982), Marciano (1984), and Maeda (1988) discussed possibilities of including these multidimensional theories into the cosmological scenario of the expanding Universe and found that the low-energy limits to the fundamental constants might vary over the cosmological time. Variations of the coupling constants of strong and electroweak interactions might then cause the masses of elementary particles to change. Note that an increase of the proton mass by 0.08\\% would lead to transformation of protons into neutrons (by electron capture), resulting in destruction of atoms in the Universe. As demonstrated by Kolb, Perry, \\& Walker (1987) and Barrow (1987), observational bounds on the time evolution of extra spatial dimensions in the Kaluza--Klein and superstring theories can be obtained from limits on possible variations of the coupling constants. Damour \\& Polyakov (1994) have developed a modern version of the string theory which assumes cosmological variations of the coupling constants and hadron-to-electron mass ratios. Therefore the parameters of the theory can be restricted by testing cosmological changes of these ratios.  The present value of the proton-to-electron mass ratio is $\\mu=1836.1526645\\,(57)$ (CODATA, 1997). Obviously, any significant variation of this parameter over a small time interval is excluded, but such variation over the cosmological time $\\sim 1.5\\times 10^{10}$ yr remains a possibility.  This possibility can be checked by analyzing spectra of high-redshift QSOs.  The first analysis of this kind has been performed by Pagel (1977), who obtained a restriction $|\\dot{\\mu}/\\mu|< 5\\times 10^{-11}{\\rm ~yr}^{-1}$ on the variation rate of $\\mu$ by comparison of wavelengths of H\\,I and heavy-ion absorption lines, as proposed by Thompson (1975). This technique, however, could not provide a fully conclusive result, since the heavy elements and hydrogen usually belong to different interstellar clouds, moving with different radial velocities. In this paper we use another technique, based on an analysis of H$_2$ absorption lines only.  One object suitable for such analysis is the $z = 2.811$ absorption system toward PKS 0528--250, in which Levshakov \\& Varshalovich (1985) identified molecular hydrogen absorption lines based on a spectrum obtained by Morton et al.\\ (1980).  Foltz, Chaffee, \\& Black (1988) have presented a limit to possible variation of $\\mu$ based on their observations of PKS 0528--250.  Their analysis did not, however, take into account wavelength-to-mass sensitivity coefficients, hence their result appeared to be not well grounded. Subsequently the spectrum of Foltz, Chaffee, \\& Black (1988) was reappraised by Varshalovich \\& Levshakov (1993), who obtained $|\\Delta\\mu/\\mu| < 0.005$ at the redshift $z=2.811$, and by Varshalovich \\& Potekhin (1995), who obtained $|\\Delta\\mu/\\mu| < 0.002$ at the $2\\sigma$ significance level.  (Here $\\Delta\\mu/\\mu$ is the fractional variation of $\\mu$.) More recently, Cowie \\& Songaila (1995) used a new spectrum of PKS 0528$-$250 obtained with the Keck telescope to arrive at the 95\\% confidence interval $-5.5 \\times 10^{-4} < \\Delta\\mu/\\mu < 7 \\times10^{-4}$.  Here we present a profile fitting analysis of a new, high-resolution spectrum of PKS 0528$-$250, obtained in November 1991 with the Cerro-Tololo Inter-American Observatory (CTIO) 4 m telescope. We have calculated the wavelength-to-mass sensitivity coefficients for a larger number of spectral lines and employed them in the analysis, which yields the strongest observational constraint yet to possible $\\mu$ variation over the cosmological time scale (eq.~[\\ref{limits}] below). ",
        "conclusions": "We have obtained a constraint to the variation rate of the proton-to-electron mass ratio $\\mu$. Two fitting procedures have been used, one of which simultaneously takes into account all observed spectral regions and transitions, while the other is applied to each spectral feature separately. The two techniques, applied to two different sets of spectral intervals, have resulted in similar upper bounds on $\\Delta\\mu/\\mu$, at the level $\\,\\sim2\\times10^{-4}$. The obtained restriction on $\\dot{\\mu}/\\mu$ (\\ref{constraint}) is by an order of magnitude more stringent than the limit set previously by Varshalovich \\& Potekhin (1995), who used a spectrum with a lower spectral resolution. Moreover, it is much more restrictive than the estimate of Cowie \\& Songaila (1995), based on high-resolution Keck telescope observations. There are two reasons for the higher accuracy of the present estimate. First, our fitting procedure simultaneously takes into account all observed spectral regions and transitions.  This is particularly important because many of the transitions are blended, even at the spectral resolution of the spectrum used by Cowie \\& Songaila (1995). A separate analysis of spectral lines leads to larger statistical errors, as we have shown explicitly in Section 4.2. Second, we include a larger number of transitions between excited states of the H$_2$ molecule (83 spectral lines, compared with 19 lines used by Cowie \\& Songaila), many of which have higher wavelength-to-mass sensitivity coefficients $K_i$. The larger interval of $K_i$ values results in a higher sensitivity to possible mass ratio deviations.  The method used here to determine the variation rate of $\\mu$ could be formally less sensitive than the one based on an analysis of relative abundances of chemical elements produced in the primordial nucleosynthesis (Kolb et al., 1986). However, the latter method is very indirect since it depends on a physical model which includes a number of additional assumptions. Therefore the present method seems to be more reliable.  Quite recently, Wiklind \\& Combes (1997) used a similar method (following Varshalovich \\& Potekhin, 1996) in order to infer limits on time variability of masses of molecules CO, HCN, HNC and the molecular ion HCO$^+$ from high-resolution radio observations of rotational lines in spectra of a few low-redshift ($z<1$) quasars. The result reported in this paper constrains the mass of the H$_2$ molecule, and thus the proton mass, at much larger $z$. These constraints may be used for checking the multidimensional cosmological models which predict time-dependences of fundamental physical constants. The described method of the calculation of the sensitivity coefficients can also be used for analyzing any other high-redshift molecular clouds, which may be found in future observations."
    },
    "9804/hep-ph9804291_arXiv.txt": {
        "abstract": "Mixed dark matter scenario can reconcile the COBE data and the observed large scale structure.  So far the massive neutrino with a mass of a few eV has been the only discussed candidate for the hot dark matter component.  We point out that the hadronic axion in the so-called hadronic axion window, $f_{a} \\sim 10^{6}$~GeV, is a perfect candidate as hot dark matter within the mixed dark matter scenario. The current limits on the hadronic axion are summarized.  The most promising methods to verify the hadronic axion in this window are the resonant absorption of almost-monochromatic solar axions from M1 transition of the thermally excited $^{57}$Fe in the Sun, and the observation of the ``axion burst'' in water \\v{C}erenkov detectors from another supernova. ",
        "introduction": "The cold dark matter (CDM) dominated universe with scale-invariant primordial density fluctuation has been the standard theory of structure formation.  After COBE has found the finite density fluctuation in the cosmic microwave background radiation (CMBR), the standard CDM scenario was found to give too much power on smaller scales.  Many modifications to the standard CDM scenario were proposed which solve the discrepancy: by introducing a small Hot Dark Matter (HDM) component~\\cite{aph9707285}, by ``tilting'' the primordial density fluctuation spectrum~\\cite{tilt}, by assuming a finite cosmological constant~\\cite{Lambda-CDM}, or by introducing particles (such as $\\nu_{\\tau}$) whose decay changes the time of radiation-matter equality~\\cite{cdm+mnu}.  At this point, there is no clear winner among these possibilities.\\footnote{However, a large ``tilt'' is difficult to obtain in many inflationary models. $\\tau$CDM can be tested well by $B$-factory experiments in the near future~\\cite{hph9709411}.  The recent data from high-redshift supernovae prefer $\\Lambda$CDM \\cite{Saul}, but the possible evolution of supernovae needs to be excluded by more systematic comparison between nearby and high-$z$ supernovae.}  In this letter, we revisit the mixed dark matter (MDM) scenario from the particle physics point of view.  This scenario has attracted strong interests because there has been a natural candidate for the HDM component: massive neutrino(s).  A neutrino with a mass of a few eV can naturally contribute to a significant fraction of the current universe.  However, it has not been easy to incorporate the HDM together with other neutrino ``anomalies,'' unless all three generation neutrinos (possibly together with a sterile neutrino) are almost degenerate, and their small mass splittings explain various ``anomalies.''  Such a scenario may be viewed as fine-tuned. Especially, the atmospheric neutrino anomaly is quite significant statistically now thanks to the SuperKamiokande experiment, which suggests the mass squared difference of $\\Delta m^{2} = 10^{-3} - 10^{-2}~{\\rm eV}^{2}$ between the muon and tau neutrinos.  If we view the situation from the familiar hierarchical fermion mass matrices, it suggests the tau neutrino mass of 0.03 -- 0.1~eV, and it appears difficult to accommodate the HDM based on massive neutrinos.  We point out that the hadronic axion~\\cite{KSVZ} can be an alternative motivated candidate for the HDM component in the MDM model.  Axion has been proposed as a solution to the strong CP problem in the QCD, and the hadronic axion (or KSVZ axion) is one version which predicts small coupling of the axion to the electron.  There has been known a window of $f_{a} \\sim 10^{6}$~GeV allowed by existent astrophysical and cosmological constraints if the axion coupling to photons is suppressed accidentally.  This is referred to as the ``hadronic axion window.'' Our main observation is that this window gives exactly the right mass of $m_{a} \\sim \\mbox{a few eV}$ and the number density of the axion appropriate for the HDM component in the MDM scenario. ",
        "conclusions": "In this letter, we have pointed out that the hadronic axion in the hadronic axion window ($f_a\\sim 10^6{\\rm ~GeV}$) can automatically be a good candidate of the Hot Dark Matter component in the mixed dark matter scenario.  In order to evade an astrophysical constraint from the background UV light, axion-photon-photon coupling has to be suppressed in the hadronic axion window, probably by an accidental cancellation. This scenario may be tested by detecting the axion burst from a future supernova in water \\v{C}erenkov detectors, or detecting solar axions using resonant absorption."
    },
    "9804/astro-ph9804071_arXiv.txt": {
        "abstract": "We present CCD BVI photometry of the old open cluster Berkeley 21, one of the most distant clusters in the Galactic anticentre direction, and possibly the lowest metallicity object in the open clusters sample. Its position and metal abundance make it very important for the study of the Galactic disc.  Using the synthetic Colour - Magnitude Diagram method, we estimate values for distance modulus \\mmm = 13.4--13.6,  reddening \\ebv = 0.74--0.78 (with possible differential absorption), and age =  2.2--2.5 Gyr. ",
        "introduction": "Old open clusters cover a large range of distances, metallicities, and ages (Friel 1995), and that warrants their use in investigations of the chemical and dynamical evolution of our Galaxy. To study  the metallicity and age distribution of open clusters with Galactocentric distance, and avoid unnecessary and dangerous biases, a key requisite is homogenous analysis of very accurate observational data, as discussed by, e.g, Janes \\& Phelps (1994, JP94) Carraro \\& Chiosi (1994, CC94), Friel (1995), Twarog et al. (1997,TAAT97).  This is the fifth paper of a series dedicated to the examination of old open clusters of different ages and metallicities, and located at different Galactic radii: for them we measure in a homogenous way distance, age, reddening and metallicity. These quantities are derived from comparison of the observed colour-magnitude diagrams (CMDs) to synthetic ones generated by a numerical code based on stellar evolution tracks and taking into account theoretical and observational uncertainties (Tosi et al. 1991). These simulations are much more powerful than the classical isochrone fitting method to study the evolutionary status of the analysed region and have been successfully applied both to nearby irregular galaxies (Greggio et al. 1998 and references therein) and to galactic open clusters (NGC2243: Bonifazi et al. 1990; Cr261: Gozzoli et al. 1996; NGC6253:  Bragaglia et al. 1997; NGC2506:  Marconi et al. 1997).  Berkeley 21 (Be21) is located toward the Galactic anticentre, at coordinates RA(1950) = 5:48:42, DEC(1950) = 21:46, and l$_{\\rm II}$ =187$^{\\circ}$, b$_{\\rm II}$ = $-2.5^{\\circ}$. It has already been observed by Christian \\& Janes (1979, hereafter CJ), but their photographic CMD is very shallow, barely reaching the main sequence Turn-Off (TO). They deduced a substantial reddening (\\ebv $\\simeq$ 1.0), a large distance modulus (\\mmm $\\simeq$ 16), and a quite young age ($\\sim 10^8$ yr). Much better data have been presented by Phelps et al. (1994, PJM94) in their compilation of old open clusters, providing \\mmm = 13.9$\\pm$0.2 and an age of 2.8 Gyr, derived on the basis of $\\delta V$=1.6 ($\\delta V$ being the magnitude difference between TO and clump stars, JP94). The metallicity has been estimated by medium-resolution spectroscopy (Friel \\& Janes 1993, FJ93), but its actual value strongly depends on the adopted reddening (\\ebv = 0.7$\\pm$0.2, Janes 1991), with [Fe/H]= $-0.97^{+0.3}_{-0.1}$ dex. This large uncertainty, given the fact that Be21 defines the lowest metallicity limit of the open clusters sample and is one of the clusters most distant from the Galactic centre, is a further limitation for studies of the (possible) age and distance relations with chemical abundance in the Galactic disc (see also Twarog et al. 1997).  In Section 2 we describe the observations and data analysis; in Section 3 we present the derived CMDs involving BVI photometry and discuss the presence of binary stars. In Section 4 we compare observed and synthetic CMDs and derive metallicity, age, distance and reddening. Finally, conclusions will be reviewed in Section 5.  \\begin{table} \\begin{center} \\caption{Log of the observations. The cluster field has its centre at RA(2000) = 5:51:46, DEC(2000) = +21:48:45. The off-cluster field has coordinates: RA(2000) = 5:52:08, DEC(2000) = +21:53:53} \\begin{tabular}{lcccccccc} \\hline\\hline \\multicolumn{1}{c}{Night} &\\multicolumn{1}{c}{Field} &\\multicolumn{1}{c}{Tel.} &\\multicolumn{3}{c}{Exposure in seconds} &\\multicolumn{1}{c}{Seeing} \\\\ & & & B & V & I & excursion \\\\ \\hline Mar 4 &Centre &Danish & 120 & 120 & 120 & 0.90-1.20\\\\ Mar 4 &Centre &Danish &  -  & 120 & 900 & 0.85-1.00\\\\ Mar 4 &Centre &Danish &  -  & 600 & - & 0.95-1.10\\\\ Mar 5 &Centre &Danish & 120 & 120 & 120 & 1.00-1.10\\\\ Mar 5 &Centre &Danish &1500 & 900 &  -  & 0.95-1.10\\\\ Mar 14 &Ext.  &Dutch & 120  & 60  &  60 & 1.10-1.30\\\\ Mar 14 &Ext.  &Dutch & 1200 & 480 & 480 & 1.20-1.30\\\\ Mar 14 &Ext.  &Dutch & 1200 & 480 & 480 & 1.15-1.30\\\\ \\hline \\end{tabular} \\end{center} \\label{tab-log} \\end{table}  \\begin{figure*} \\vspace{14cm} \\special{psfile=figmap.ps vscale=95 hscale=95 voffset=-175 hoffset=-25} \\caption{Map of the observed field, taken from our $V,B-V$ photometry. North is up and East left.} \\label{fig-map} \\end{figure*} ",
        "conclusions": "Although the fits of observed and simulated diagrams are not as satisfying for Be21 as they have been for other systems examined with the same method, we have been able to determine a fairly consistent confidence interval for its distance, age, reddening and metallicity (see Table 5) by selecting the most reliable among the models described in the previous section. They place it among the old metal poor open clusters, in a region far from the Galactic centre and of moderately high reddening.  \\subsection{Distance and reddening}  We have derived a distance slightly smaller than previous studies, while our evaluation of the reddening is fairly consistent with past works. No previous indication of differential reddening was given in the literature, but our data definitely show it.  JP94 found, for the red clumps of 23 open clusters with \\dv$\\ge$1.0 (i.e. older than about 1.5 Gyr) a  mean absolute magnitude $M_V=0.9 \\pm 0.4$, and a mean intrinsic colour  $(B-V) = 0.95 \\pm 0.10$. In our case, these mean values, when applied to the observed $V$=16.80 and $B-V$=1.55, would imply $(m-M)_V$=15.90, and \\ebv=0.60, or $(m-M)_0$=14.04. From our best simulations, we obtain instead $(m-M)_0 \\simeq $ 13.5, \\ebv=0.76, corresponding to $(m-M)_V$=15.86. In other words, the clump-based distances seem to agree, but the colour of the clump stars seems to be quite different from the mean. Part of this discrepancy may be due to the high reddening affecting Be21, whose differential effect on blue and red stars leads to an apparent shrink by 0.04 of the true colour difference between TO and clump stars (Fernie 1963, Twarog 1998 private communication). On the other side, TAAT97, derive a  mean  $M_V=0.6 \\pm 0.1$ for ten clusters not too metal-rich ranging from NGC7789 to Mel66, i.e. approximately from 1 to 5 Gyr, for which they try to measure the distance in a quite homogenous way. This translates, in our case, to $(m-M)_V$=16.2; since they do not cite a mean intrinsic $(B-V)$ no further comparison with our best choice for the reddening is possible.  There are no completely reliable reddening determination for this cluster since the UBV data of CJ do not reach the MS, but our determination and that by Janes (1991) agree well. We have further compared our finding with what is expected from the spatial distribution of interstellar extinction near the Galactic plane. To this end we have considered the studies of FitzGerald (1968, fig. 3h) and Neckel \\& Klare (1980, fig. 6i). In both cases, a reasonable estimate deduced from their  data for low Galactic latitudes and the right longitude, is \\ebv $\\sim$ 0.8; FitzGerald's (1968) observations also allow for a lower value, closer to 0.5, but seem to exclude the high values, close to 1, needed by the lower metallicity tracks of any group.  Janes (1991) and JPM94  give a distance modulus \\mmm=13.9$\\pm$0.2,  somewhat larger than our results for every set of tracks. Carraro et al. (1998), working on the same data, cite a Galactocentric distance of 14.5 Kpc, also implying a distance modulus \\mmm$\\simeq$13.9. We have no good explanation for this difference, but we must emphasize that adopting \\mmm=13.9 we would be forced to select younger ages, and this would have two major drawbacks: a worse disagreement with literature ages (see next), and a worse reproduction of the MS shape in the synthetic CMDs.  \\subsection{Age}  Also in the case of the age, we seem to have found a value lower than given in literature. We can explain the discrepancy partly  by the different techniques adopted, partly by the better quality of our data. The various parameter combinations all converge to a fairly small range of possible ages (2.1 to 2.8 Gyr, with favorite age around 2.2 Gyr). In fact, the only largely discrepant value found in our analysis is for the FRANEC98 Z=0.01 tracks, a value in strong disagreement with the spectroscopically determined metallicity. Despite the uncertainties involved in the age determination with our method, we consider it still more reliable than ages derived by other means.  Nonetheless, it is not always feasible to determine the age of a cluster with the proper method of synthetic diagram fitting: to do so, high quality data, both deep and precise, are needed, and the process itself is complex. To apply this technique to all the objects of interest takes a long time, while the properties of the whole sample of open clusters are needed to study the Disc population and evolution. For this reason, several parameterizations of cluster ages,  based on a handful of well studied objects, have been proposed. These methods, if uncertain in absolute value for the single cluster, yield a reasonable age ranking for the cluster system.  These parameterizations are usually based on a difference, in magnitude and/or in colour, between well recognizable points of the CMD (usually the TO and the red clump), as this is much easier to measure than any absolute quantity. Note though that the precise definition of the two points, and especially of the TO, changes among authors. We will cite here the three following examples: i) Anthony-Twarog \\& Twarog (1985 and later works) use the magnitude difference between the red giant clump and {\\it the brightest point at the TO} ($\\delta V_T$) coupled with the difference in colour between the red giant branch at the position of the clump and the bluest point of the TO ($\\delta(B-V)_T$); ii) JP94 use a similar \\dv, but measured between the red clump and {\\it the inflection point between the MS TO and the base of the giant branch}; iii) CC94 define their $\\Delta V$ as JP94, but assume that {\\it the reference TO luminosity is 0.25 mag fainter than observed in the CMD}, to take into account the fact that presence of unresolved binaries tends to brighten the TO region. In the case of Be21, we have: \\dv=1.8 (our measure), $\\delta V_T$=1.2, \\dv(JP94)=1.6, $\\Delta V$=1.55. All these different definitions try to circumvent the problem that the TO point is not always easily identified in open clusters, due to field stars and binaries contamination and/or paucity of stars. The strength of our kind of analysis is that we do not judge on the basis of the observed CMD alone: we know from the tracks the exact location of the TO in each of our simulated CMDs. Given this, we have chosen to measure the magnitude difference as defined above at what we believe to be the true TO, i.e. at the point corresponding to the hottest MS point in the evolutionary tracks.  JP94 correlated \\dv ~with cluster ages. The calibration of their Morphological Age Index (MAI, expressed in Gyr) translates for Be21 to an age of 2.8 Gyr (based upon \\dv=1.6, PJM94), marginally inconsistent with what we get from the direct comparison with evolutionary tracks (see Section 4). The age difference does not arise from any discrepancy in the two sets of data: we find  \\dv=1.8, quite consistent with the value given by PJM94 considering that we measure it in a slightly different way. However, we have found in the past (e.g. in the case of NGC2506, Marconi et al. 1997) that the MAI tends to overestimate ages.  Anthony-Twarog \\& Twarog (1985, revised by Twarog \\& Anthony-Twarog 1989) proposed the so called Morphological Age Ratio (MAR), defined as MAR = $\\delta V_T / \\delta (B-V)_T$ (see above). This index is independent of reddening and almost independent of metallicity (Anthony-Twarog \\& Twarog 1985, Buonanno et al. 1989). The calibration of the relation between MAR and ages has changed from: age = 1.4$\\times$MAR Gyr (Anthony-Twarog \\& Twarog 1985) to: age = 2.0$\\times$MAR Gyr (Twarog \\& Anthony-Twarog 1989). Applying their definition to our CMD, we find the values given in Table 6, and an age of about 2.7 to 3.9 Gyr, depending on the adopted calibration. With no attempt to give a new calibration of the MAR-age relation, simply adopting for the clusters we studied the parameters and ages in Table 6, we find:  age = 2.3 $\\times$ MAR -- 2.6 ($r.m.s.$ = 0.9). We did not include Be17 in this computation: it represents an extreme of the interpolation and we felt that the parameters derived from the published diagrams were too insecure. This relation gives for Be21 an age of 1.9 Gyr.  Carraro et al. (1998) derive for Be21 an age of 3.1 Gyr, based on the PJM94 data and the synthetic CMD method using the Padova tracks. The difference with our results, obtained employing the same sets of tracks (although we do not interpolate in metallicity as they do), may perhaps be explained simply with the worse quality of the observational data they use. Certainly, in no case are we able to reach self-consistently such a large age.  \\subsection{Metallicity}  We note that our method is unable to solve the problem of the cluster precise metal abundance. The comparison with evolutionary tracks can only give a coarse indication of metallicity. Too many variables are present in tracks computation to discriminate metallicity to such an extent. In fact, tracks nominally closer to the metallicity derived for Be21 from spectroscopic measurements ([Fe/H]=--0.97, or Z $\\simeq$ 0.002, FJ93) appear less consistent with the observed CMD than tracks more metal rich, because of the large colour extent of the subgiant branch and, in some cases, of an excessively high reddening required to reproduce the observed colour of the MS.  Anyway, what can be said is that the best fits are obtained for the slightly more metal-rich combinations, i.e. for Z=0.006 (\\fe $\\simeq$ -0.5) or 0.004 (\\fe $\\simeq$ -0.7) as compared to 0.001 (\\fe $\\simeq$ -1.3). This would go in favour of a metal abundance slightly higher than measured by FJ93.  There is the possibility that the FJ93 scale may be underestimating cluster metallicities.  TAAT97 compared it  with abundances based on DDO photometry and found the FJ93 values systematically low. Another example may be the couple of clusters examined by Gratton \\& Contarini (1994): they observed two giants in each cluster at high-resolution and high S/N (R=30,000, S/N $\\simeq$ 100) and found for NGC2243 and Mel66 the values [Fe/H]=--0.48 and --0.38 respectively, to be compared with [Fe/H]=--0.56 and --0.51 (FJ93).  FJ93 emphasized the fact that the actual value derived for the cluster metallicity from their spectra is strongly dependent on the adopted reddening: the $\\pm$0.2 mag error on reddening in Janes (1991) allows for a formal uncertainty of $\\pm$0.3 dex in metallicity. They also find marginal support for a \\ebv ~value on the higher side, hence for a metal abundance slightly higher than the [Fe/H]=--0.97 they give. This goes in the same direction suggested by our comparisons, even if we do not find any convincing evidence for a larger reddening. Finally, we have identified the four stars studied by FJ93 (N$_{CJ}$ = 50, 406, 413, 415a, which correspond to N$_{our}$ = 50, 67, 20, 51 respectively), to check if they may be influenced by the differential reddening we found; but we consider it quite improbable, since all the four objects are within 1 arcmin from the centre of the cluster.  No conclusive word can be said on Be21 metallicity, which would instead be important to know with high precision, since it could represent the lowest value for the open clusters in our Galaxy. A decisive answer would come from high resolution spectroscopy coupled with fine abundance analysis on the four stars examined by FJ93, already known to be cluster members.  \\bigskip\\bigskip\\noindent ACKNOWLEDGEMENTS  \\noindent We warmly thank A. Chieffi, M. Limongi and, specially, O. Straniero for having not only distributed the new FRANEC tracks in advance of publication, but even in format suitable for our purposes. We also thank J.C. Mermilliod for kindly making available his invaluable BDA open clusters database and for useful comments. We are grateful to the referee (Bruce Twarog) for his comments, extremely useful both to improve the clarity of the paper and for future applications. The bulk of the numerical code for CMD simulations has been provided by Laura Greggio. This research has made use of the Simbad database, operated at CDS, Strasbourg, France."
    },
    "9804/astro-ph9804137_arXiv.txt": {
        "abstract": "We analyze the data of low--energy cosmic--ray $\\bar p$ spectrum, recently published by  the BESS Collaboration, in terms of newly calculated fluxes for secondary antiprotons and for a possible contribution of an exotic signal due to neutralino annihilation in the galactic halo. We single out the relevant supersymmetric configurations and discuss their explorability with experiments of direct search for particle dark matter and at accelerators. We discuss how future measurements with the Alpha Magnetic Spectrometer (AMS) on the Shuttle flight may disentangle the possible neutralino--induced contribution from the secondary one. ",
        "introduction": "A recent analysis \\cite{bess95} of the data collected by the balloon--borne BESS spectrometer on cosmic--ray antiprotons during its flight in 1995 (hereafter referred to as BESS95 data) has provided the most detailed information on the low--energy cosmic--ray $\\bar{p}$'s spectrum currently available: 43 antiprotons have been detected, grouped in 5 narrow energy windows over the total kinetic--energy range $180 ~ {\\rm MeV} \\leq  T_{\\bar{p}} \\leq 1.4 ~{\\rm GeV}$. With this experiment the total number of measured cosmic--ray antiprotons in balloon--borne detectors over a period of more than 20 years \\cite{all,hof,mitchell,moiseev,barbiellini} has more than doubled. Most remarkably, the BESS95 data provide a very useful information over the low--energy part of the ${\\bar p}$ flux, where a possible distortion of the spectrum expected for secondary $\\bar{p}$'s (i.e., antiprotons created by interactions of primary cosmic--ray nuclei with the interstellar medium) may reveal the existence of cosmic--ray antiprotons of exotic origin (for instance, due to pair annihilation of relic particles in the galactic halo \\cite{th,noi,mitsui}, to evaporation of primordial black holes \\cite{mitsui,kww} or to cosmic strings \\cite{strings}). In fact, a possible discrimination between primary (exotic) and secondary $\\bar p$'s is based on the different features of their low--energy spectra: in this energy regime ($ T_{\\bar p} \\lsim $ 1 GeV)  interstellar (IS) secondary $\\bar p$ spectrum is expected to drop off very markedly because of kinematical reasons \\cite{gaisser}, while exotic antiprotons show a milder fall off. However, as will be discussed later on, this discrimination  power is somewhat hindered by solar modulation and by some other effects affecting particle diffusion in the Galaxy.  In Fig. 1 we report the cosmic--ray $\\bar p$ flux at the top of the atmosphere (hereafter referred to as TOA flux) measured by BESS95 \\cite{bess95}. For experimental data referring to other measurements with much less statistics see Refs.\\cite{all,hof,mitchell,moiseev,barbiellini}. Also displayed in Fig. 1 are the minimal, median and maximal fluxes expected for secondary antiprotons at the time of the BESS95 data taking. These fluxes have been derived with a procedure which is described in detail in Secs. II--V.  A comparison of the BESS95 data with the theoretically expected fluxes for secondary $\\bar p$'s, as displayed in Fig. 1, leads to the following considerations: i) the experimental data are consistent with the theoretically expected secondary flux, within the experimental errors and the theoretical uncertainties; however, ii) the experimental flux seems to be suggestive of a flatter behaviour, as compared to the one expected for secondaries ${\\bar p}$'s. Thus, natural questions arise, such as: a) how much room for exotic ${\\bar p}$'s would there be in the BESS95 data, for instance in case the secondary flux is approximately given by the median estimate of Fig. 1, b) how consistent with the current theoretical models would be the interpretation of the BESS95 data in terms of a fractional presence of exotic antiprotons, and  c) how this interpretation might be checked by means of independent experiments? In the present note we address these questions within an interpretation of a possible excess of $\\bar p$'s at low energies in terms of primary antiprotons generated by relic neutralinos in the galactic halo \\cite{note1}.  The present analysis \\cite{note2} is mostly meant to a clarification of many theoretical points which will be even more crucial, when a much more statistically significant experimental information on low--energy cosmic--ray antiprotons will be made available by forthcoming experiments: AMS on the precursor Shuttle flight in May 1998 and on the International Space Station Alpha (ISSA) in January 2002 \\cite{ams}, the satellite--borne PAMELA experiment \\cite{pamela} and balloon--borne measurements \\cite{spill}.  Our paper is organized as follows. In Sec. II we discuss the cosmic--ray IS proton spectrum which will be subsequently employed in deriving the secondary antiprotons. In the same section we also illustrate how we treat the solar modulation to connect the IS spectra to the corresponding TOA fluxes. In Sec. III we discuss the sources of cosmic antiprotons, both of primary and of secondary origins. Cosmic rays diffusion properties are derived in Sec. IV; the TOA $\\bar p$ spectra are given in Sec. V. In Sec. VI we compare our theoretical fluxes with the BESS95 data and single out the neutralino configurations which may be relevant for the present problem. Secs. VII and VIII are devoted to an analysis on how these supersymmetric configurations  can be explored by direct searches for relic neutralinos and by experimental investigation at accelerators. Conclusions and perspectives in terms of the forthcoming measurements of low--energy cosmic--ray $\\bar p$'s  are illustrated in Sec. IX. ",
        "conclusions": "We have presented a new  analysis of the cosmic--ray antiprotons flux, expected on the basis of secondary $\\bar p$'s, generated by interactions of cosmic--ray primaries with the interstellar medium, and of a possible exotic primary source of $\\bar p$'s, originated by neutralino--neutralino annihilations in the Galactic halo.  Improvements over previous calculations of secondaries depend mostly on: i) the use of a two--zone propagation model for diffusion of cosmic rays in the halo instead of the standard leaky box model; ii) the inclusion of an energy--loss effect in the propagation properties of cosmic rays (important for the antiproton low energy range considered in this paper); iii) the use of the new data on primary cosmic--ray proton spectrum, as measured by IMAX  \\cite{imaxp} and CAPRICE \\cite{caprice}.  The neutralino--induced $\\bar p$ flux has been evaluated in a MSSM at the electroweak scale, which incorporates all current accelerator constraints. Use of supergravity--inspired unification conditions  at large energy scale has been avoided in order not to arbitrarily constrain the neutralino phenomenology \\cite{bere}.  Solar modulation of the antiproton flux has been improved by analyzing the most complete set of data over the solar cycles \\cite{papini} and the data on the proton spectrum of Refs. \\cite{imaxp,caprice}.   We have found that the most statistically relevant data on cosmic--ray antiprotons at low--energy \\cite{bess95} leave some room for a possible signal from neutralino annihilation in the galactic halo. We have discussed how the relevant supersymmetric configurations may be explored with direct experiments for particle dark matter search and at accelerators. We have shown how the interplay between measurements of cosmic--ray $\\bar p$'s and direct search experiments for relic particles is very intriguing and quite important in view of the significant improvements expected in these two classes of experiments in the near future. The present analysis stresses the great interest for the forthcoming AMS measurements with the Shuttle flight and on the ISSA \\cite{ams}, as well as for other future measurements with balloon--borne experiments (IMAX\\cite{mitchell}, BESS\\cite{moiseev}) and with satellites (PAMELA) \\cite{pamela}, for disentangling the secondary $\\bar p$ flux from a possible primary signal of exotic nature. As an example, we give in Fig. 18 the distribution of measurements expected for AMS with the Shuttle flight according to two different hypothesis: a) dominance of the secondary contribution (lower sequence  of crosses), b) significant contribution due a neutralino--induced signal (upper sequence of crosses). In our evaluation of the expected measurements we have taken into account geomagnetic cut--off effects and the expected AMS overall acceptance \\cite{batt}.   \\begin{center} {\\bf Acknowledgments} \\end{center}  We thank Roberto Battiston and Aldo Morselli for  interesting discussions about  experimental aspects related to the present paper. P.S. would like to thank the French Programme National de Cosmologie for its support.   \\newpage \\begin{table} \\caption{Values of the parameters in the expressions (\\ref{eq:energy}) and (\\ref{eq:rigidity}) for the IS proton flux and of the solar--modulation parameter $\\Delta$. These values are obtained by best--fitting the data of Refs.[19-20] with Eqs. (\\ref{eq:energy}) and (\\ref{eq:rigidity}), either over the entire energy range or only over the high--energy ($T_p \\geq 20$ GeV) range. First and third sets of values refer to 3--parameters fits (with Eqs.(\\ref{eq:energy}) and (\\ref{eq:rigidity}), respectively), second and fourth sets refer to 2--parameters fits at fixed $\\Delta$, (with Eqs.(\\ref{eq:energy}) and (\\ref{eq:rigidity}), respectively). } \\begin{center} \\begin{tabular}{|c|c|c|c|}   \\hline  &     IMAX          &    CAPRICE     &  Comments \\\\ \\hline \\hline  A         & 12,300$\\pm$1,700  & 17,600$\\pm$500 & entire energy\\\\ $\\alpha$  & 2.67$\\pm$0.03     & 2.81$\\pm$0.01  &   range    \\\\ $\\Delta$  & 510$\\pm$40        & 390$\\pm$ 5     &            \\\\ \\hline  A         & 12,300$\\pm$3,000  & 19,600$\\pm$ 3,000 &    \\\\ $\\alpha$  & 2.67$\\pm$0.06     & 2.85$\\pm$0.04     & $T_p \\geq$ 20 GeV  \\\\ $\\Delta$  & 510 (fixed)       & 390 (fixed)       &        \\\\ \\hline   B         & 16,200$\\pm$2,000  & 26,000$\\pm$1,200 & entire energy\\\\ $\\gamma$  & 2.73$\\pm$0.03     & 2.91$\\pm$0.02    &   range    \\\\ $\\Delta$  & 795$\\pm$35        & 710$\\pm$10       &         \\\\ \\hline  B         & 13,700$\\pm$4,100  & 22,800$\\pm$ 3,700 &      \\\\ $\\gamma$  & 2.69$\\pm$0.06     & 2.88$\\pm$0.04     &  $T_p \\geq$ 20 GeV  \\\\ $\\Delta$  & 795 (fixed)       & 710 (fixed)       &     \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\vfill \\eject  \\newpage  \\begin{center} {\\bf Figure Captions} \\end{center}   {\\bf Fig. 1} - TOA antiproton flux as a function of the antiproton kinetic energy. The experimental points are the BESS95 data \\cite{bess95}. The curves are the median (solid line), minimal (dotted line) and maximal (dashed line) secondary fluxes calculated in this paper, solar--modulated at the time of the BESS95 measurement.  {\\bf Fig. 2} - TOA spectra of IMAX (full circles) \\cite{imaxp} and of CAPRICE (open circles) \\cite{caprice} with our best--fit curves with parametrization of Eq. (\\ref{eq:energy}). (The error bars are not shown when they are smaller than the dimension of the circles.)   {\\bf Fig. 3} - TOA spectra of IMAX (full circles) \\cite{imaxp} and of CAPRICE (open circles) \\cite{caprice}. The solid, dotted and dashed lines denote the median, minimal and maximal IS proton fluxes, respectively, as discussed in Sec. II. (The error bars are not shown when they are smaller than the dimension of the circles.)   {\\bf Fig. 4} - The grammage $\\Lambda_e$ of the CNO primary elements (solid) as inferred from a two--zone diffusion model of the propagation of cosmic rays in the Galaxy. It is plotted as a function of the kinetic energy per nucleon. The dashed curve features the grammage corresponding to protons while the dotted lines delineate the interval of escape lengths inferred from the Ficenec {\\it et al.} \\cite{Ficenec} observations on $^{3}$He at TOA energies comprised between 100 MeV/n and 1.6 GeV/n. \\label{fig:grammage}  {\\bf Fig. 5} - IS secondary antiproton spectra as functions of the ${\\bar p}$ kinetic energy. Solid, dotted and dashed lines denote the fluxes obtained from the median, minimal and maximal IS primary proton fluxes. The dot--dashed line denotes the median ${\\bar p}$ spectrum, when the ${\\bar p}$ energy losses are neglected.  {\\bf Fig. 6} - Coefficient $C_{\\rm susy}(T_{\\bar p},f)$ as a function of the ${\\bar p}$ kinetic energy for different values of the flattening parameter $f$.  {\\bf Fig. 7} - Time variation of the solar--modulation parameter $\\Delta$. Full circles represent the best--fit values to the PGS average fluxes at minima (MIN) and at maxima (MAX) and to the fluxes  of IMAX \\cite{imaxp} and of CAPRICE \\cite{caprice}; the open circle refers to the BESS95 data taking period and the cross denotes the extrapolated  value of $\\Delta$ at the time relevant for the future AMS Shuttle flight (May 1998).  {\\bf Fig. 8} - Solar modulation of the IS median secondary antiproton flux calculated in this paper. Solid line is the IS spectrum; dashed (dotted) line is the solar--modulated spectrum at minima (maxima).  {\\bf Fig. 9} - Solar modulation of the IS antiproton flux, due to neutralino annihilation for the representative neutralino configuration with $m_{\\chi} = 62$ GeV, $P = 0.98$ and $\\Omega_{\\chi} h^2 = 0.11$. Solid line is the IS spectrum; dashed (dotted) line is the solar--modulated spectrum at minima (maxima).   {\\bf Fig. 10} - TOA antiproton fluxes versus the antiproton kinetic energy. The BESS95 data \\cite{bess95} are shown by crosses. The dashed line denotes the median secondary flux, the dotted one denotes the primary flux due to neutralino annihilation in the halo for a neutralino configuration with $m_{\\chi} = 62$ GeV, $P = 0.98$ and $\\Omega_{\\chi} h^2 = 0.11$. Solid line denotes the calculated total flux.  {\\bf Fig. 11} - Scatter plots for configurations of set $M$ (a) and set $N$ (b) in the P--$\\tan \\beta$ plane.  {\\bf Fig. 12} - Scatter plots for configurations of set $M$ (a) and set $N$ (b) in the P--$m_{\\chi}$ plane.  {\\bf Fig. 13} - Scatter plots for configurations of set $R$ in the P--$\\tan \\beta$ plane (a) and in the P--$m_{\\chi}$ plane (b).  {\\bf Fig. 14} - Scatter plots for configurations of set $R$ in the P--$\\tan \\beta$ plane (a) and in the P--$m_{\\chi}$ plane (b) for a flattening of $f = 0.5$.  {\\bf Fig. 15} -  Scatter plot of the values of $\\xi \\sigma_{\\rm scalar}^{\\rm nucleon}$ versus the neutralino mass for the configurations of set $M$ (a) and of set $N$ (b). The open curve denotes the (90 \\% C.L.) upper bound obtained from experimental data of Ref. \\cite{few}. The region delimited by a closed contour is the one singled out by the experiment of Ref. \\cite{dama} as possibly indicative of an annual modulation effect. The total local dark matter density is normalized here to the value $\\rho_l = 0.4$ GeV cm$^{-3}$. The dashed line shows the discovery potential in case of an improvement by a factor of 10 in current sensitivities for experiments of direct search for particle dark matter.  {\\bf Fig. 16} - Correlation between $\\xi \\sigma_{\\rm scalar}^{\\rm nucleon}$ and the neutralino relic density $\\Omega_{\\chi} h^2$ for configurations of set $M$.  {\\bf Fig. 17} - Scatter plot for configurations of set $M$ in the $m_h$--$\\tan \\beta$ plane. The region on the left of the dashed line denoted by (a) is excluded by current LEP experimental data \\cite{lep182}, the one on the right of the dashed line (b) is theoretically disallowed. The other lines display the LEP reach at luminosity $L = 200$ pb$^{-1}$ and various energies \\cite{alt}: (A) discovery potential at ${\\sqrt s} = 192$ GeV; (B) discovery potential at ${\\sqrt s} = 200$ GeV; (C) exclusion at ${\\sqrt s} = 200$ GeV.   {\\bf  Fig. 18} - Expected distribution of measurements with the AMS Shuttle flight according to two different hypothesis: a) dominance of the secondary contribution (lower sequence  of crosses), b) significant contribution due a neutralino--induced signal (upper sequence of crosses). The dashed line denotes the secondary flux, the dotted one denotes the primary flux due to neutralino annihilation in the halo for a neutralino configuration with the representative values: $m_{\\chi} = 62$ GeV, $P = 0.98$ and $\\Omega_{\\chi} h^2 = 0.11$. Solid line denotes the calculated total flux."
    },
    "9804/astro-ph9804301_arXiv.txt": {
        "abstract": "We present a fully covariant and gauge-invariant calculation of the evolution of anisotropies in the Cosmic Microwave Background (CMB) radiation. We use the physically appealing covariant approach to cosmological perturbations, which ensures that all variables are gauge-invariant and have a clear physical interpretation. We derive the complete set of frame-independent linearised equations describing the (Boltzmann) evolution of anisotropy and inhomogeneity in an almost Friedmann-Robertson-Walker (FRW) Cold Dark Matter (CDM) universe. These equations include the contributions of scalar, vector and tensor modes in a unified manner. Frame-independent equations for scalar and tensor perturbations, which are valid for any value of the background curvature, are obtained straightforwardly from the complete set of equations. We discuss the scalar equations in detail, including the integral solution and relation with the line of sight approach, analytic solutions in the early radiation dominated era, and the numerical solution in the standard CDM model. Our results confirm those obtained by other groups, who have worked carefully with non-covariant methods in specific gauges, but are derived here in a completely transparent fashion. ",
        "introduction": "The cosmic microwave background radiation (CMB) occupies a central role in modern cosmology. It provides us with a unique record of conditions along our past lightcone back to the epoch of decoupling (last scattering), when the optical depth to Thomson scattering rises suddenly due to Hydrogen recombination. Accurate observations of the CMB anisotropy should allow us to distinguish between models of structure formation and, in the case of non-seeded models, to infer the spectrum of initial perturbations in the early universe. Essential to this programme is the accurate and reliable calculation of the anisotropy predicted in viable cosmological models.  Such calculations have a long history, beginning with Sachs \\&\\ Wolfe (1967) who investigated the anisotropy on large scales $(\\gtrsim 1^{\\circ})$ by calculating the redshift back to last scattering along null geodesics in a perturbed universe. On smaller angular scales one must address the detailed local processes occurring in the electron/baryon plasma prior to recombination, and the effects of non-instantaneous last scattering. These processes, which give rise to a wealth of structure in the CMB power spectrum on intermediate scales and damping on small scales (see, for example, Silk (1967, 1968)), are best addressed by following the photon distribution function directly from an early epoch in the history of the universe to the current point of observation. This requires a numerical integration of the Boltzmann equation, and has been carried out by many groups, of which Peebles \\&\\ Yu (1970), Bond \\&\\ Efstathiou (1984, 1987), Hu \\&\\ Sugiyama (1995), Ma \\&\\ Bertshcinger (1995), Seljak \\&\\ Zaldarriaga (1996) is a representative sample.  The calculation of CMB anisotropies is simple in principle, but in reality is plagued with subtle gauge issues (Stoeger, Ellis, \\&\\ Schmidt 1991; Stoeger et al.\\ 1995; Challinor \\&\\ Lasenby 1998). These problems arise because of the gauge-freedom in specifying a map $\\Phi$ between the real universe (denoted by $S$) and the unperturbed background model(denoted by $\\bar{S}$)~\\cite{ellis89a}, which is usually taken to be a Friedmann-Robertson-Walker (FRW) universe. The map $\\Phi$ identifies points in the real universe with points in the background model, thus defining the perturbation in any quantity of interest. For example, for the density $\\rho$ as measured by some physically defined observer, the perturbation at $x\\in S$ is defined to be $\\delta \\rho(x)\\equiv \\rho(x) - \\bar{\\rho}(\\bar{x})$, where $\\bar{\\rho}$ is the equivalent density in the background model, and $x$ maps to $\\bar{x}$ under $\\Phi$. The map $\\Phi$ is usually (partially) specified by imposing coordinate conditions in $S$ and $\\bar{S}$. Any residual freedom in the map $\\Phi$ after the imposition of the coordinate conditions (gauge-fixing) gives rise to the following gauge problems: (i) the map cannot be reconstructed from observations in $S$ alone, so that quantities such as the density perturbation, which depend on the specific map $\\Phi$, are necessarily not observable; (ii) if the residual gauge freedom allows points in $\\bar{S}$ to be mapped to physically inequivalent points in $\\bar{S}$ in the limit that $S=\\bar{S}$, then unphysical gauge mode solutions to the linearised perturbation equations will exist.  There are several ways to deal with the gauge problems described above. In the earliest approach~\\cite{lifshitz46}, one retains the residual gauge freedom (in the synchronous gauge) but keeps track of it so that gauge mode solutions can be eliminated. Furthermore, the final results of such a calculation must be expressed in terms of the physically relevant, gauge-invariant quantities. Although there is nothing fundamentally wrong with this approach if carried out correctly, it suffers from a long history littered with confusion and errors. The need to express results in terms of gauge-invariant variables suggests that it might be beneficial to employ such variables all along as the dynamical degrees of freedom in the calculation. A further advantage of such an approach is that gauge modes are automatically eliminated from the perturbation equations when expressed in terms of gauge-invariant variables. This is the approach adopted by Bardeen (1980), who showed how to construct gauge-invariant variables for scalar, vector and tensor modes in linearised perturbation theory, by taking suitable linear combinations of the gauge-dependent perturbations in the metric and matter variables. This approach has been used in several calculations of the CMB anisotropy (see, for example, Abbott \\&\\ Schaefer (1986) and Panek (1986)). However, the Bardeen variables are not entirely satisfactory. The approach is inherently linear, so that the variables are only defined for small departures from FRW symmetry. Furthermore, the approach assumes a non-local decomposition of the perturbations into scalar, vector and tensor modes at the outset, each of which is then treated independently. As a result, the Bardeen variables are only gauge-invariant for the restricted class of gauge-transformations that respect the scalar, vector and tensor splitting. Finally, although the Bardeen variables are gauge-invariant, they are not physically transparent, in that, in a general gauge, they do not characterise the perturbations in a manner that is amenable to simple physical interpretation.  An alternative scheme for the gauge-invariant treatment of cosmological perturbations was given by Ellis \\&\\ Bruni (1989) (see also, Ellis, Hwang, \\&\\ Bruni (1989)) who built upon earlier work by Hawking (1966). In this approach, which is derived from the covariant approach to cosmology/hydrodynamics of Ehlers and Ellis (Ehlers 1993; Ellis 1971), the perturbations are described by gauge-invariant variables that are covariantly defined in the real universe. This ensures that the variables have simple physical interpretations in terms of the inhomogeneity and anisotropy of the universe. Since the definition of the covariant variables does not assume any linearisation, exact equations can be found for their evolution, which can then be linearised around the chosen background model. Furthermore, the covariant approach does not employ the non-local decomposition into scalar, vector or tensor modes, at a fundamental level. If required, the decomposition can be performed at a late stage in the calculation to aid solving the equations. Even if one denies that working with gauge-invariant variables is a significant advantage, the key advantage of the covariant approach, however, is that one is able to work exclusively with physically relevant quantities, satisfying equations that make manifest their physical consequences.  The covariant and gauge-invariant approach has already been applied to the line of sight calculation of CMB anisotropies under the instantaneous recombination approximation (Dunsby 1997; Challinor \\&\\ Lasenby 1998), and has been used to obtain model-independent limits on the inhomogeneity and anisotropy from measurements of the CMB anisotropy on large scales~\\cite{maartens95}. In this paper, we extend the methodology developed in these earlier papers, to give a full kinetic theory calculation of CMB anisotropies valid on all angular scales. Our motivation for reconsidering this problem is two-fold. Firstly, it is our belief that the covariant and gauge-invariant description of cosmological perturbations provides a powerful set of tools for the formulation of the basic perturbation equations, and their subsequent interpretation, which are superior to the techniques usually employed in such calculations for the reasons discussed above. In particular, by applying covariant methods for the problem of the generation of CMB anisotropies, we can expect the same advantages of physical clarity and unification that have already been demonstrated in other areas, (Ellis et al.\\ 1989; Bruni, Ellis, \\&\\ Dunsby 1992; Dunsby, Bruni, \\&\\ Ellis 1992; Dunsby, Bassett, \\&\\ Ellis 1996; Tsagas \\&\\ Barrow 1997). The approach described here brings the underlying physics to the fore, and can only help to consolidate our rapidly growing understanding of the physics of CMB anisotropies. Furthermore, although we only consider the linearised calculation here, the extension of these methods to the full non-linear case is quite straightforward (Maartens, Gebbie, \\&\\ Ellis 1998). Our second motivation is to perform an independent verification of the results of other groups (for example, Ma \\&\\ Bertschinger (1995)), with a methodology that is free from any of the gauge ambiguities that have caused problems and confusion in the past. Given the potential impact on cosmology of the next generation of CMB data, we believe that the above comments provide ample justification for reconsidering this problem.  For definiteness we consider the cold dark matter (CDM) model, although the methods we describe are straightforward to extend to other models. We have endeavoured to make this paper reasonably self-contained, so we begin with a brief overview of the covariant approach to cosmology and define the key variables we use to characterise the perturbations in Section~\\ref{sec_cov}. We then go on to present a complete set of frame-independent equations describing the evolution of the matter components and radiation in Section~\\ref{sec_eqs} in an almost-FRW universe (with arbitrary spatial curvature). These equations, which employ only covariantly defined, gauge-invariant variables, are independent of any harmonic analysis; they describe scalar, vector and tensor perturbations in a unified manner. Many of the equations have simple Newtonian analogues, and their physical consequences are far more transparent than the equations that underlie the metric-based approaches. Equations pertinent to scalar modes, see Section~\\ref{sec_scal}, and tensor modes, see Section~\\ref{sec_tens}, can be obtained from the full set of equations with very little effort, and are useful at this late stage in the calculation as an aid to solving the linearised equations. A significant feature of this approach is that a covariant angular decomposition of the distribution functions is made early on in the calculation, before any splitting into scalar, vector and tensor modes. This allows scalar, vector and tensor modes to be treated in a more unified manner. In particular, the azimuthal dependence of the moments of the distribution functions does not have to be put in by hand (after inspection of the azimuthal dependence of the other terms in the Boltzmann equation), as happens in most metric-based calculations. This is particularly significant for tensor modes where the required azimuthal dependence is non-trivial and is different for the two polarisations of gravitational waves. We consider the equations for scalar modes in considerable detail. We present the integral solution of the Boltzmann multipole equations in a $K=0$ almost-FRW universe, and discuss the relation between line of sight methods (which employ lightlike integrations along the lightcone) and the Boltzmann multipole approach (where a timelike integration is performed). We derive analytic solutions for scalar modes in the early radiation dominated universe, that are used as initial conditions for the numerical solution of the scalar equations, the results of which we describe in Section~\\ref{sec_num}. In Section~\\ref{sec_tens} we give a brief discussion of the tensor equations in the covariant approach. The covariant angular decomposition naturally gives rise to a set of variables that describe the temperature anisotropy in a more direct manner than in the conventional metric-based approaches. This is particularly apparent for tensor perturbations, where the CMB power spectrum at a given multipole $l$ is determined by the $l-2$, $l$ and $(l+2)$-th moments of the conventional decomposition of the photon distribution function, which obscures the physical interpretation of these moments. Finally, we end with our conclusions in Section~\\ref{sec_conc}. Ultimately, our results confirm those of other groups (for example, Ma \\&\\ Bertschinger (1995)) who have performed similar calculations by working carefully in specific gauges, but are obtained here with a unified methodology that is more physically transparent and less prone to lead to confusion over subtle gauge effects.  We employ standard general relativity and use a $(+---)$ metric signature. Our conventions for the Riemann and Ricci tensors are fixed by $[\\nabla_{a},\\nabla_{b}] u^{c} = -{\\clr_{abd}}^{c} u^{d}$, and $\\clr_{ab} \\equiv {\\clr_{acb}}^{c}$. Round brackets around indices denote symmetrisation on the indices enclosed, and square brackets denote antisymmetrisation. We use units with $c=G=1$ throughout, and a unit of distance of $\\mpc$ for numerical work. ",
        "conclusions": "\\label{sec_conc}  We have shown how the full kinetic-theory calculation of the evolution of CMB anisotropies and density inhomogeneities can be performed in the covariant approach to cosmology (Ehlers 1993; Ellis 1971), using the gauge-invariant perturbation theory of Ellis \\&\\ Bruni (1989). Adopting covariantly-defined, gauge-invariant variables throughout ensured that our discussion avoided the gauge ambiguities that appear in certain gauges, and that all variables had a clear, physical interpretation. We presented a unified set of equations describing the evolution of photon and neutrino anisotropies and cosmological perturbations in the CDM model, which were independent of a decomposition into scalar, vector or tensor modes and the associated harmonic analysis. Although we only considered the case of linear perturbations around an FRW universe here, it is straightforward to extend the approach to include non-linear effects (Maartens et al. 1998), which should allow a physically transparent discussion of second-order effects on the CMB. Indeed, the ease with which one can write down the exact equations for the physically relevant variables is one of the major strengths of the covariant approach.  The linear equations describing scalar modes and tensor modes were obtained from the full set of equations in a straightforward and unified manner, highlighting the advantage of having the full equations (independent of the decomposition into scalar, vector and tensor modes) available. For the scalar case, the Boltzmann multipole equations for the moments of the distribution functions obtained here were equivalent to those usually seen in the literature. However, for tensor modes, the covariant approach led naturally to a set of moment variables that more conveniently describe the temperature anisotropy than those usually employed. For scalar modes, we discussed the solution of the perturbation equations in detail, including the integral solution of the Boltzmann multipole equations and the relation between the timelike integrations performed in the multipole approach to calculating CMB anisotropies, and the lightlike integrations of the line of sight approach. The numerical solution of the scalar equations in a $K=0$, almost-FRW, CDM universe were also discussed. Our numerical results provide independent confirmation of those of other groups, (see, for example, Ma \\&\\ Bertschinger (1995) and Seljak \\&\\ Zaldarriaga (1996)), who have obtained their results by employing non-covariant methods in specific gauges. Typically, these methods require one to keep careful track of all residual gauge-freedom, both to enable identification of any gauge-mode solutions, and to ensure that the final results quoted are gauge-invariant (and hence observable). Fortunately, the isotropy of the photon distribution function in an exact FRW universe ensures that the CMB power spectrum, as calculated from the gauge-dependent perturbation to the distribution function, is gauge-invariant for $l\\geq 1$.  We hope to have shown the ease with which the covariant approach to cosmology can be applied to the problem of calculating CMB anisotropies. The covariant and gauge-invariant method discussed here frees one from the gauge problems that have caused confusion in the past, and focuses attention on the physically relevant variables in the problem and the underlying physics. Future work in this area will include the discussion of non-linear effects (Maartens at al. 1998), the inclusion of polarisation, and the effects of hot dark matter, all of which can be expected to bring the same advantages of physical clarity and transparency that we hope to have demonstrated here."
    },
    "9804/astro-ph9804184_arXiv.txt": {
        "abstract": "This paper describes the first optical spectroscopic survey of class I sources (also known as embedded sources and protostars) in the Taurus-Auriga dark cloud.  We detect 10 of the 24 known class I sources in the cloud at 5500--9000 \\AA.  All detected class I sources have strong H$\\alpha$ emission; most also have strong [O~I] and [S~II] emission.  These data -- together with high quality optical spectra of T Tauri stars in the Taurus-Auriga cloud -- demonstrate that forbidden emission lines are stronger and more common in class I sources than in T Tauri stars.  Our results also provide a clear discriminant in the frequency of forbidden line emission between weak-emission and classical T Tauri stars. In addition to strong emission lines, three class I sources have prominent TiO absorption bands.  The M-type central stars of these sources mingle with optically visible T Tauri stars in the HR diagram and lie somewhat below both the birthline for spherical accretion and the deuterium burning sequence for disc accretion. ",
        "introduction": "Examining the earliest phases of low mass stellar evolution requires observations of protostars deeply embedded in the dense cores of nearby molecular clouds.  These ``class I'' sources (\\cite{lad87}) have blackbody-like spectral energy distributions that peak at wavelengths of 30--100 \\mum~and bolometric luminosities, $L_{\\rm b} \\approx$ 0.1--100 \\lsun~(\\cite{ada87}; \\cite{mye87}; \\cite{wil89}; \\cite{ken90}; \\cite{gre94}).  Despite many detailed studies of their circumstellar environments (see, for example, \\cite{ta91a},b; \\cite{and93}; \\cite{ter93}; \\cite{mor92}, 1995; \\cite{bon96}; \\cite{gom97}; \\cite{hog97}; \\cite{whi97}), understanding the stellar physics of these objects has proved elusive. Comparisons of observed bolometric luminosity functions with models is straightforward but controversial (\\cite{wil89}; \\cite{ken90}; \\cite{fl94a}, 1994b).  The apparent lack of photospheric features in many objects has led several groups to abandon the HR diagram as a means for testing stellar evolutionary tracks of the youngest stars. These groups have proposed the bolometric temperature (\\cite{mye98} and references therein), the submillimeter flux (\\cite{sar96}), and the visual extinction (\\cite{ada90}) to replace effective temperature and have developed models to place evolving pre--main-sequence stars in their modified evolutionary diagrams.  The accuracy of these techniques remains uncertain, because the methods are new and incompletely tested.  In this paper, we report an optical spectroscopic survey designed to detect photospheric absorption features from the central stars of class I sources in the Taurus-Auriga cloud.  Although the line-of-sight extinction to the central star is large, $A_V \\approx$ 30--60 mag (\\cite{whi97}), large ground-based telescopes can detect optical light scattered off cavities in the infalling envelopes of many objects. Optical data also provide the best measure of spectral types for pre--main-sequence stars.  In general, I-band and J-band data are least contaminated by emission from an accretion disc and its associated boundary layer or magnetic accretion column (\\cite{kh90}). However, the very large continuum veiling detected on near-IR spectra of some class I sources (\\cite{cas92}; \\cite{gr96a}, 1996b) favors I-band spectra, because the J-band veiling can be large if the disc extends to the stellar photosphere (\\cite{kh90}; \\cite{ken96}).  Finally, optical spectra of class I sources allow an unambiguous comparison with optically brighter T Tauri stars, which have known spectral types in a well-calibrated system (see, for example, Kenyon \\& Hartmann 1995; KH95 hereafter).  Our results provide the first optical detection of M-type absorption features in an embedded protostar.  We identify TiO absorption bands in three Taurus-Auriga class I sources; one other star may have TiO features and a fifth star may have K-type absorption features.  We use optical spectra of T Tauri stars to calibrate the spectral types of class I sources and then construct a complete HR diagram for the Taurus-Auriga cloud.  These data, coupled with new evolutionary tracks for protostars accreting from discs and two spectral types derived from near-IR spectra (\\cite{gr96b}), show that class I sources in Taurus-Auriga mingle with T Tauri stars and lie below the birthline in the HR diagram.  We also detect strong emission lines on the spectra of all protostars. Forbidden emission from [N~II] and [S~II] is much more common among class I sources than older, optically brighter stars having the same bolometric luminosity.  The fluxes of forbidden emission lines also seem stronger among class I sources than other pre--main-sequence stars in the cloud.  We find no evidence that the permitted emission lines, such as H$\\alpha$ and He~I, are more common or stronger than in T Tauri stars.  These results extend and confirm previous conclusions that jet activity declines as a pre--main-sequence star contracts to the main-sequence.  We describe our observations in Sec. 2, explain our results in Sec. 3, and conclude with a brief discussion in Sec. 4. ",
        "conclusions": "In the previous sections, we have described the first optical spectroscopic survey of class I, embedded sources in a single molecular cloud.  We supplemented these data with high quality optical spectra of a representative sample of older and optically brighter T Tauri stars.  The combined set of spectra show that the optical spectra of class I sources qualitatively resemble the optical spectra of T Tauri stars.  Our analysis further reveals common physical properties and substantial differences between class I sources and T Tauri stars, as summarized below.  Our data provide the first indication that the distribution of stellar spectral types among class I sources may not be very different from that of WTTS and CTTS.  Of the five class I sources with strong optical continua, one (L1489 IRS) is a continuum + emission source, three are M-type stars, and another (04264+2433) may have an M-type central star.  To the best of our knowledge, {\\it these are the first low mass protostars with measured optical spectral types.} This sample is too small for a meaningful comparison with the distribution of spectral types among more evolved pre--main-sequence stars in the cloud.  We note, however, that the median spectral type for WTTS and CTTS is K7-M0 and that the frequency of continuum + emission sources is $\\sim$ 5\\%--10\\% (KH95).  Published observations indicate other similarities between class I sources and older pre--main-sequence stars in Taurus-Auriga. First, class I sources have the same intrinsic near-IR colors as do CTTS.  Whitney \\etal (1997) show that the observed near-IR colors of class I sources can be modeled as a CTTS surrounded by an infalling envelope with an optical extinction, $A_V \\approx$ 30--60 mag. This analysis leads to the conclusion that the radiation from the star and inner disc of a class I source is similar to that of a T Tauri star (see also \\cite{gr96b}; Calvet \\etal 1997 reach a different conclusion). Second, the bolometric luminosity distributions of class I sources, CTTS, and WTTS are indistinguishable (KH95).  All three groups of pre--main-sequence stars have median luminosities of $L_{\\rm b} \\approx$ 0.5--0.8 \\lsun.  This unexpected result is supported by the positions of class I sources in the HR diagram.  Our data show that class I sources have luminosities and effective temperatures very similar to those of CTTS and WTTS in the cloud. These conclusions are surprising, because a class I source should have a larger luminosity once it has accreted nearly all of its final mass, and this luminosity should decline with time as the star approaches the main sequence (see, for example, \\cite{sta83}, 1988; \\cite{pal93}; \\cite{har97}; Fig. 8).  The current sample, however, is too small to test stellar models in detail. The errors in luminosity and effective temperature are also too large. Observations with the next generation of ground-based telescopes will undoubtedly expand the sample, reduce the errors, and provide better tests of protostellar accretion theory.  One feature that distinguishes class I sources is their strong forbidden-line emission.  As a group, class I sources are much more likely to have forbidden-line emission than CTTS or WTTS (Fig. 7). This result confirms previous conclusions from imaging data (e.g., \\cite{gom97}) and indicates that class I sources are more likely to drive outflows than CTTS or WTTS (see also \\cite{bon96}; \\cite{mor92}, 1994).  The equivalent widths of the forbidden lines are also larger in class I sources than in CTTS or WTTS.  Although some large equivalent widths may be due to very weak optical continua, the [S~II] equivalent widths in HH31 IRS2 -- a class I source with a prominent TiO absorption band -- are larger than observed in {\\it any} CTTS in our sample (see Tables 1--2).  Deeper optical spectra of our sample and other class I sources would clarify this point.  Our sample is not large enough to test whether class I sources also have more prominent {\\it permitted} emission lines than CTTS. The median H$\\alpha$ equivalent width for class I sources, $\\sim$ 90 \\AA, is much larger than the median equivalent width for CTTS, $\\sim$ 30--40 \\AA.  This difference is roughly what we expect if class I sources have larger continuum veiling than CTTS (\\cite{cas96}; \\cite{gre97}) and if the H$\\alpha$ equivalent width correlates with veiling (\\cite{har95} and references therein).  However, the frequency of He~I $\\lambda\\lambda$5876, 6678 emission among class I sources is roughly comparable to that among CTTS. We measure a He~I emission frequency of 50\\% among 6 class I sources with reasonable signal-to-noise at 6000 \\AA, 57\\% among 14 flat-spectrum sources, and 65\\% among 46 class II sources. For both emission lines, the class I sample is probably biased against small equivalent widths, because class I sources without emission lines are probably fainter than sources with emission lines. A deeper survey with a larger telescope could enlarge the sample of class I sources with high quality optical spectra.  These data would provide a good test for differences in the distribution of H$\\alpha$ equivalent widths between class I sources and CTTS.  These results fit into the general picture of Taurus-Auriga class I sources developed in KH95 and in Kenyon \\etal (1990).  In this picture, class I sources are envelopes of gas and dust falling into the central star-disc system at rates of a few $\\times~10^{-6}~\\msunyr$ (see also \\cite{ada87}; \\cite{ke93a}, 1993b; \\cite{whi97}). Bell \\& Lin (1994) show that the stable accretion rate through the disc onto the central star is either very low -- $\\lesssim$ a few $\\times~10^{-7}~\\msunyr$ -- or very high -- $\\gtrsim$ a few $\\times~10^{-5}~\\msunyr$ -- compared to the infall rate.  The disc spends most of its time in the low accretion rate state; the disc mass then  slowly increases with time until it reaches a critical level and evolves to the high accretion rate state. This model explains the low observed luminosities of nearly all class I sources as well as the occasional high luminosity of a source such as L1551 IRS5.  Models with time-dependent disc accretion also qualitatively account for the evolution of forbidden and permitted emission lines in pre--main-sequence stars.  We expect the time-averaged accretion rate through the disc to decline as the envelope disperses.  If the H$\\alpha$ and other permitted emission lines of class I sources form in the accretion region of the inner disc as in CTTS, then the median H$\\alpha$ equivalent width should decline as a pre--main sequence star evolves from a class I source into a CTTS and then into a WTTS.  Most models for jet formation link the mass loss rate in the jet to the mass accretion rate through the disc (see, for example, \\cite{cab90}; \\cite{naj94}; \\cite{sh94a}, 1994b), so we expect forbidden emission to decline as well. Explaining the observations of emission line equivalent widths with a quantitative model of a dispersing envelope and evolving disc, however, is currently beyond our reach.  Finally, our results further demonstrate the advantages of optical spectra.  Recent surveys of larger samples of class I sources using near-IR spectroscopy have yielded only two spectral types each in Taurus-Auriga (\\cite{cas96}; \\cite{gr96b}) and $\\rho$ Oph (\\cite{gr96b}, 1997). Casali \\& Eiroa (1996; see also \\cite{cas92}; \\cite{gr96b}, 1997) conclude that continuum emission from dust in a circumstellar disc or envelope veils photospheric absorption features on near-IR spectra of class I sources.  Preliminary results further suggest that this veiling is larger in class I sources than in CTTS or WTTS (\\cite{cas96}; \\cite{gr96b}, 1997). Dust emission is much weaker relative to a normal stellar photosphere at shorter wavelengths, $\\lesssim 1~\\mu$m, so optical spectra may yet provide the best measure of spectral types in class I sources.  \\vskip 6ex We thank the staffs of the MMT, Palomar, and Whipple Observatories for assistance with our observations.  Fred Chaffee kindly acquired several spectra of the class II sources listed in Table 1. Susan Tokarz reduced the FAST spectra and graciously assisted with the reduction of the MMT and Palomar spectra.  We also thank M. Geller, M. G\\'omez, C. Lada, A. Mahdavi, and B. Whitney for advice and comments.  The suggestions of an anonymous referee improved our presentation. Observations at the Palomar Observatory were made as part of a continuing collaborative agreement between Palomar Observatory and the Jet Propulsion Laboratory. Portions of this research were supported by the National Aeronautics and Space Administration through grant NAGW-2919 and by the Space Telescope Science Institute through grant GO-06132.01-94A. C.A.T. thanks the Royal Society and the Hungarian Academy of Sciences for an exchange fellowship during the majority of his contribution to this project.  \\vfill \\eject"
    },
    "9804/astro-ph9804229_arXiv.txt": {
        "abstract": "We present and discuss 25 spectra obtained in November 1996, covering all phases of the CAL 87 binary system.  These spectra are superior both in signal-to-noise and wavelength coverage to previously published data so that additional spectral features can be measured.  Photometry obtained on the same nights is used to confirm the ephemeris and to compare with light curves from previous years.  Analysis of the color variation through the orbital cycle has been carried out using archival MACHO data.\\footnote{This paper utilizes public domain data obtained by the MACHO Project, jointly funded by the US Department of Energy through Lawrence Livermore National Laboratory under contract W7405-Eng-48, the National Science Foundation through the Center for Particle Astrophysics of the University of California undercooperative agreement AST-8809616, and the Mount Stromlo and Sidings Springs Observatory by the Bilateral Science Technology and Regional Development.}  When a barely resolved red field star is accounted for, there is no ($V-R$)-color variation, even through eclipse.  There have been substantial changes in the depth of minimum light since 1988; it has decreased more than 0.5 mag in the last several years.  The spectral features and radial velocities are also found to vary not only through the 0.44-day orbit but also over timescales of a year or more.  Possible interpretations of these long-term changes are discussed.  The 1996 spectra contain phase-modulated Balmer absorption lines not previously seen, apparently arising in gas flowing from the region of the compact star.  The changes in emission-line strengths with orbital phase indicate there are azimuthal variations in the accretion disk structures.  Radial velocities of several lines give different amplitudes and phasing, making determination of the stellar masses difficult.  All solutions for the stellar masses indicate that the companion star is considerably less massive than the degenerate star.  The Balmer absorption-line velocities correspond to masses of $\\sim$1.4M$_{\\odot}$ for the degenerate star and $\\sim$0.4M$_{\\odot}$ for the mass donor.  However, the strong He II emission lines indicate a much more massive accreting star, with M$_X>$4M$_{\\odot}$. ",
        "introduction": "CAL 87 has long been known as one of only a small number of luminous X-ray binaries in the Large Magellanic Cloud (Long, Helfand, \\& Grabelsky 1981, Pakull et al.\\ 1988).  Its optical spectrum, with He II and H emission lines on a very blue continuum, shows it to be similar to galactic low-mass X-ray binaries.  However, its X-ray spectrum reveals CAL 87 to be one of the rare, very luminous (L$_{bol}\\geq10^{38}$ erg s$^{-1}$) supersoft sources (SSS) which have little or no radiation above $\\sim$0.5 keV (e.g. Tr\\\"umper et al.\\ 1991, Greiner 1996).  The SSS are widely thought to be binaries in which a white dwarf is undergoing rapid accretion from a more massive companion, leading to steady nuclear burning on the surface of the white dwarf (van den Heuvel et al.\\ 1992).  CAL 87 is unique among supersoft sources in having both optical and X-ray eclipses (Callanan et al.\\ 1989, Cowley et al.\\ 1990: CSCH) which provide extra information about the disk structure and in principle help to constrain the stellar masses.  In the original spectroscopic data of CSCH, He II 4686\\AA\\ emission was shown to move with K$=$40 km s$^{-1}$ and proper phasing with respect to the eclipse so that it was interpreted as due to orbital motion of the compact star.  CSCH concluded that the compact star had a mass $\\ge$6M$_{\\odot}$, and hence these data implied the presence of a black hole.  However, optical spectra taken a few years later, combined with velocities from lines in the far UV (Hutchings et al.\\ 1995), showed a quite different behavior.  The velocity amplitude was larger and the phasing was very different, indicating that at times the velocities are not entirely due to orbital motion and the line-formation regions change.  One problem is that spectroscopic determination of the masses is not entirely straightforward since the nearly edge-on view of the system ($i\\sim78^{\\circ}$) causes complications in the line profiles due to motions in the accretion disk.  In this paper we report on a series of spectroscopic data with improved signal-to-noise (S/N) and new, concurrent photometry which we obtained at CTIO in 1996 in order to conduct a more thorough investigation of CAL 87. ",
        "conclusions": "We have discovered that the photometric eclipse depth has changed over recent years, and this variation may even be cyclic.  The increase of light observed at mid-eclipse may be due to a geometrical change in the disk structure, without a significant change in total luminosity -- such as a bright structure above or below the disk plane.  Alternatively, the disk may grow in size in its plane, accompanied by extra absorption due to the outflowing material, along lines-of-sight near the center.  The emergence of the Balmer absorption in the 1996 spectra suggests the latter scenario.  In 1994 November we obtained three optical spectra of CAL 87 which were reported by Hutchings et al.\\ (1995) in their discussion of eight ultraviolet spectra observed by HST in 1995 January.  We have remeasured the optical spectra, although they are of considerably lower quality than the new data and only show the strongest lines.  In 1994--5 He II 4686\\AA\\ and 1640\\AA\\ emission-line velocities show a very different phasing (maximum velocity near $\\phi\\sim0.9$) and higher amplitude (K$\\sim$70 km s$^{-1}$) compared to both the CSCH data (K=40 km s$^{-1}$) and present data (K$\\sim$30 km s$^{-1}$).  The equivalent width of 4686\\AA\\ was lower in the 1994 spectra than in 1996 by a factor of nearly two.  In two of the three 1994 optical spectra weak absorption features are visible at H$\\beta$ and H$\\gamma$, similar to those seen in 1996, but the low S/N make the velocities unreliable.  Thus, it appears that there are long-term spectral changes that may include significant non-orbital motions.  The O VI lines might give a clean measure of orbital motion, as they arise in the inner disk, but they are very weak, resulting in a large velocity scatter.  The measured velocity amplitude (K=35 km s$^{-1}$) leads to the same conclusion as CSCH, that the compact star must be massive, with M$_X>$4M$_{\\odot}$.  If instead the velocity amplitude of the compact star is shown by the Balmer absorption lines (K=73 km s$^{-1}$), then the resulting mass diagram is almost identical to that of the supersoft binary SMC 13 (Crampton et al.\\ 1997) and lower masses are determined.  We have argued above that the measured He II velocity amplitude may be somewhat lower than the actual velocity of the compact star, because the phasing of absorptions will decrease the apparent velocity extremes, while non-orbital motions within the system add at other phases. We point out that the He II velocities in CSCH, when the eclipse was deeper, gave K=40  km s$^{-1}$. The Balmer-line velocities are also likely to contain some non-orbital motions.  Thus K=73 km s$^{-1}$ is a reasonable upper limit for the motion of the degenerate star.  Figure~\\ref{mass} shows the resulting masses corresponding to these two extreme K values.  The masses are constrained by the fairly well-known inclination ($i\\sim70^{\\circ}-80^{\\circ}$) and by the requirement that the mass-losing star fills its Roche lobe. No main sequence star fills the Roche lobe defined by these plots, for orbital inclinations larger than $\\sim$30$^o$. Thus, the mass-losing star must have evolved off the main sequence to cause mass-exchange. The lowest mass that can thus evolve within a Hubble time is $\\sim$ 0.4 M$_{\\odot}$, and these are marked in the diagram at the appropriate inclination value. These therefore correspond to the lowest X-ray star masses under these assumptions, and just allow a white dwarf in the K=73 km s$^{-1}$ case. It may be possible that the mass-losing star has very low mass by having lost most of it in some earlier event, and is still filling its Roche lobe now. However, in any of these cases the X-ray star is more massive and the resulting masses are not those expected by the `standard' model (e.g. van den Heuvel et al.\\ 1992) in which it is assumed that the compact star is a $\\sim$1M$_{\\odot}$ white dwarf and the donor star has a mass of $\\sim$2.0 M$_{\\odot}$.  The spectrum in eclipse shows no signs of a late type spectrum (see Figures~\\ref{med},~\\ref{binned},~\\ref{bin2}), consistent with the absence of an evolved more massive secondary. Unfortunately, the complex changes in the disk spectrum make it difficult to be more definitive than this at present, but it appears advisable to revisit the evolutionary scenarios to accomodate a mass-losing star that is less massive than the compact star, not only in CAL 87 but also for the other supersoft X-ray binaries.  We discuss separately (Cowley et al.\\ 1998) the overall mass determinations for a number of supersoft binary systems.  Some of the supersoft X-ray binaries have been found to have highly displaced lines indicating the presence of `jets'.  These are most noticable in two low-inclination systems RX J0513$-$69 (Crampton et al.\\ 1996) and CAL 83 (Crampton et al.\\ 1987) where the displacements are several thousand kilometers per second.  In the intermediate inclination system RX J0019$+$22 the jets show a velocity of $\\sim\\pm$800 km s$^{-1}$ from the central emission line.  The jet lines move with the same phase and velocity amplitude as the central line, suggesting that the motion of He II 4686\\AA\\ emission is indeed orbital.  Since CAL 87 is viewed nearly edge-on, any line emission coming from such jets would have a very low radial velocity and thus not be separated from the central emission profile.  The asymmetry seen in the He II 4686\\AA\\ line of CAL 87 could arise from a pair of shifted lines which have larger amplitude (or possibly slightly different phase) than the central line.  This would imply that the central line motion is not entirely orbital, but without better resolution this suggestion cannot be verified."
    },
    "9804/astro-ph9804190_arXiv.txt": {
        "abstract": "We report on Westerbork 1.4 GHz radio observations of the radio counterpart to $\\gamma$-ray burst GRB~970508, between 0.80 and 138 days after this event. The 1.4 GHz light curve shows a transition from optically thick to thin emission between 39 and 54 days after the event.  We derive the slope $p$ of the spectrum of injected electrons (d$N$/d$\\gamma_{\\rm e}\\propto\\gamma_{\\rm e}^{-p}$) in two independent ways which yield values very close to $p=2.2$. This is in agreement with a relativistic dynamically near-adiabatic blast wave model whose emission is dominated by synchrotron radiation and in which a significant fraction of the electrons cool fast. ",
        "introduction": "The peak luminosities of $\\gamma$-ray bursts (GRBs) are highly super-Eddington and require relativistic outflows (Paczy\\'{n}ski 1986; Goodman 1986).  Paczy\\'{n}ski and Rhoads (1993) pointed out that radio emission is expected as a result of the interaction between such a relativistic outflow and an external medium, as is, e.g., observed in extragalactic jet sources (see also Katz 1994; M\\'{e}sz\\'{a}ros and Rees 1997). They estimated that the strongest GRBs may be followed by transient ($\\sim 10$ mJy) radio emission at intervals ranging from minutes (for a distance $d$ $\\sim$ 10$^{5}$ pc) to several weeks ($d$ $\\sim$ 10$^{9}$ pc).  However, searches for radio counterparts through follow-up observations (Frail et al. 1994,1997a; Koranyi et al. 1995; Galama et al. 1997a,b) were without success until recently.  With the rapid accurate location capability of the Wide Field Cameras (WFCs; Jager et al. 1995) onboard the Italian-Dutch X-ray observatory BeppoSAX (Piro et al. 1995) it has recently become possible to detect fading X-ray, optical and radio counterparts to GRBs (Costa et al. 1997a; Piro et al. 1997a; Groot et al. 1997a,b; Van Paradijs et al. 1997; Galama et al. 1997c,1998a; Sahu et al. 1997; Bond 1997; Metzger et al. 1997; Frail et al. 1997b,c; Bremer et al. 1998; Halpern et al. 1997).  These observations have settled the discussion on the GRB distance scale (`galactic halo' versus `cosmological', see e.g. Fishman and Meegan 1995, Lamb 1995, Paczy\\'nski 1995):  GRBs occur at Gpc distances.  GRB 970508 is the first GRB to be detected in the radio (Frail et al. 1997b,c); the radio source position coincides with that of the optical (Bond 1997) and X-ray (Piro et al. 1997b) afterglow sources. Assuming that the variations of the source at 4.86 and 8.46 GHz are due to interstellar scintillation (ISS), their damping with time is consistent with a highly relativistically expanding shell passing a diameter of $\\approx 3\\mu$as (Frail et al. 1997b). VLBI observations show that the source is unresolved ($<$ 0.3 mas, Taylor et al. 1997).  We here report on the results of 1.4 GHz radio observations of GRB 970508, made with the Westerbork Synthesis Radio Telescope (WSRT) between 0.80 and 138 days after the burst occurred. ",
        "conclusions": "}  The observed optical spectral slope $\\alpha$ and the optical power law decay of the light curve $F_{\\nu} \\propto t^{\\delta}$ is not consistent with the expected relation for the simplest blast wave model ($\\delta = 3\\alpha/2$; e.g. Wijers, Rees and M\\'esz\\'aros 1997). The observed power law decay value, $\\delta = -1.141 \\pm 0.014$ ($t > 2$ days, Galama et al. 1998b; see also Pedersen et al. 1998, Castro-Tirado et al. 1998, Sokolov et al. 1998) would imply $\\alpha = -0.761 \\pm 0.009$, while in the optical passband $\\alpha = -1.12 \\pm 0.04$ is observed (Galama et al. 1998c, from here on Paper II; see also Sokolov et al. 1998).  In the following we show that this may be explained by rapid cooling of a significant fraction of the electrons.  A population of electrons with a power-law distribution of Lorentz factors $\\gamma_{\\rm e}$ (d$N$/d$\\gamma_{\\rm e}\\propto\\gamma_{\\rm e}^{-p}$) above some minimum value $\\gamma_{\\rm m}$ emits a power law synchrotron spectrum above the frequency $\\nu_{\\rm m}$ (corresponding to radiating electrons with $\\gamma_{\\rm m}$; e.g. Rybicki \\& Lightman 1979).  Independently, above some Lorentz factor $\\gamma_{\\rm c}$ the electrons may cool rapidly, and an extra break in the spectrum is expected at the corresponding frequency $\\nu_{\\rm c}$ (Sari, Piran, \\& Narayan 1998). Beyond a certain time $t_0$ ($t_0$ is small $\\sim$ 500 sec; see Paper II) the evolution of the blast wave is adiabatic (Sari et al. 1998 and see Paper II); then $\\nu_{\\rm m} < \\nu_{\\rm c}$, and the spectrum varies as $F_{\\nu} \\propto \\nu^{-(p-1)/2}$ from $\\nu_{\\rm m}$ up to $\\nu_{\\rm c}$; above $\\nu_{\\rm c}$ it follows $F_{\\nu} \\propto \\nu^{-p/2}$ and below $\\nu_{\\rm m}$ it follows the low frequency tail, $F_{\\nu} \\propto \\nu^{1/3}$ (Sari et al. 1998). The evolution in time of the GRB afterglow is determined by the evolution of these break frequencies: $\\nu_{\\rm c} \\propto t^{-1/2}$ and $\\nu_{\\rm m} \\propto t^{-3/2}$ (both decrease with time).  The decay part of the optical R$_{\\rm c}$ (Coussins R) band light curve (in the optical passband $\\nu > \\nu_{\\rm c}$ for $t\\gsim$ 1.2 days; Paper II) goes as $F_{\\nu} \\propto t^{(2-3p)/4}$, while the spectrum is then $F_{\\nu} \\propto \\nu^{-p/2}$ (Sari et al. 1998). This allows us to make two independent measurements of $p$: using $F_{\\rm R_{\\rm c}}\\propto t^{-1.141 \\pm 0.014}$ we find $p= 2.188 \\pm 0.019$ and using $\\alpha_{\\rm opt} = -1.12 \\pm 0.04$ gives $p = 2.24 \\pm 0.08$.  The excellent agreement between the values of $p$ supports that a significant fraction of the electrons cool rapidly and that the evolution of the GRB remnant is adiabatic. Additional evidence for rapid cooling of a significant fraction of the electrons is given in Paper II.  Observations by Bremer et al. (1998) with the IRAM Plateau de Bure Interferometer (PdBI) at 86 GHz show a maximum around $\\sim$ 12 days. We identify this maximum with the break frequency $\\nu_{\\rm m}$ passing 86 GHz at $t_{\\rm m,86 GHz} \\sim$ 12 days (Paper II). We expect, the 8.46 and 1.4~GHz emission to peak at $t_{\\rm m,8.46 GHz} \\sim$ 55 days and $t_{\\rm m,1.4 GHz} \\sim$ 180 days, respectively ($\\nu_{\\rm m} \\propto t^{-3/2}$). Near day 55, a shallow maximum can be seen in the 8.46 GHz light curve (Frail et al. 1997b). Unfortunately our 1.4 GHz light curve cannot be used to test the presence of the maximum at that frequency, both due to low  signal to noise and because it ends 150 days after the burst, i.e. before the predicted maximum.   Before $\\nu_{\\rm m}$ passes 8.46 GHz at $t_{\\rm m,8.46 GHz}$ we expect the 8.46 GHz spectrum to follow the low frequency tail $F_{\\nu} \\propto \\nu^{1/3}$, while after $t_{\\rm m,8.46 GHz}$ it is expected to be $F_{\\nu} \\propto \\nu^{-(p-1)/2} = \\nu^{-0.6}$ (Sari et al. 1998 and we have used $p$ = 2.2). Thus, we predict a gradual transition between $t_{\\rm m,8.46 GHz} \\sim$ 55 days and $t_{\\rm m,4.86 GHz} \\sim$ 80 days (when also at 4.86 GHz $\\nu_{\\rm m}$ has passed) from $\\alpha$ = 1/3 to $\\alpha = -0.6$. We note that this expectation is different from blast wave models that do not include the effect of rapid cooling of a significant fraction of the electrons ($F_{\\nu} \\propto \\nu^{-1.1}$ similar to the optical slope; see e.g. Wijers et al. 1997). Also the decays at 8.46 GHz (after $t_{\\rm m,8.46 GHz} \\sim$ 55 days) and 4.86 GHz (after $t_{\\rm m,4.86 GHz} \\sim$ 80 days) are expected to be different from that in the optical and X-ray passbands, $F_{\\nu} \\propto t^{3(1-p)/4} = t^{-0.9}$; where we have used $p$ = 2.2). These predictions can be tested with the continued monitoring at the VLA at 4.86 and 8.46 GHz by Frail et al. (1998).  The radio afterglow light curves of GRB 970508 (Frail et al. 1997b and this Letter) show a much more gradual evolution than expected (see e.g. the fit to the 8.46 and 4.86 GHz data by Waxman, Kulkarni and Frail 1998). Also a constant self-absorption frequency was expected (e.g. Waxman et al. 1997) while we here show that a transition from optically thick to thin emission occurred around $\\sim$ 45 days.  For $t < t_o$ Sari et al. (1998) predict a decrease with time of the self-absorption frequency, $\\nu_{\\rm a}$, while for $t >$ $t_0$ the self-aborption frequency $\\nu_{\\rm a}$ remains constant.  The transition from optically thick to thin 1.4 GHz radiation then suggests that $t_0$ $\\sim$ 45 days. Also the 8.46 GHz light curve (Frail et al. 1997b) suggests that $t_0$ cannot be much smaller than 10 days, i.e. 10 days $\\lsim$ $t_0$ $\\lsim$ 55 days (we have extrapolated backwards in time from the 8.46 GHz peak at $t_{\\rm m}$ $\\sim$ 55 days with the expected dependence $F_{\\nu} \\propto t^{1/2}$ for times $t < t_{\\rm m}$). This is not in agreement with the finding that $t_0 \\sim 500$ sec (Paper II). However, the absence of a break in the smooth power law decay of the optical light curve from 2 to 60 days after the burst (Pedersen et al. 1998; Castro-Tirado et al. 1998; Sokolov et al. 1998; Galama et al. 1998c) shows that there is no important transition in that period. This does imply that some additional ingredient is needed; for example, Waxman et al. (1998) argue that the transition from ultrarelativistic to mildly relativistic expansion of the blast wave may explain the decrease in the self-absorption frequency $\\nu_{\\rm a}$ with time and the slow time dependence of the early radio light curves.  The excellent agreement in the derived value for $p$ ($p$ = 2.2) from the decay of the optical light curve and the optical spectral slope support an adiabatic dynamical evolution of the GRB remnant and an extra break in the synchrotron spectrum at the frequency $\\nu_{\\rm c}$ above which the radiation is from electrons which cool rapidly compared to the remnant's expansion time. We predict a transition in the radio spectral index $\\alpha_{\\rm 4.86-8.46 GHz}$ from 1/3 to --0.6, between 55 and 80 days; the light curves are predicted to decay as $F_{\\nu} \\propto t^{-0.9}$ after 55 days at 8.46 GHz and 80 days at 4.86 GHz."
    },
    "9804/astro-ph9804159_arXiv.txt": {
        "abstract": "We present observations of \\htwo\\ fluorescence at wavelengths between 1000 and 1200 \\AA\\ from the bright reflection nebula IC 63. Observations were performed with the Berkeley spectrograph aboard the \\orf-II mission \\cite{Hetal98}. To the best of our knowledge, this is the first detection of astrophysical \\htwo\\ fluorescent emission at these wavelengths (excluding planetary atmospheres). The shape of the spectrum is well described by the model of \\citeN{S89}. The absolute intensity, however, is fainter than an extrapolation from observations at longer ultraviolet wavelengths \\cite{WSBB89} by a factor of ten.  Of the mechanisms that might help reconcile these observations, optical depth effects in the fluorescing \\htwo\\ itself are the most promising (or at least the most difficult to rule out). ",
        "introduction": "In many environments, the equilibrium abundance of the hydrogen molecule (\\htwo) is determined by a balance between formation on dust grains and photodissociation pumped by photons at wavelengths below about 1108 \\AA\\ \\cite{SB82}. Photons between this threshold energy and the Lyman limit populate a large number of rovibrational states in the $B^1 \\Sigma_u^+$ and $C^1 \\Pi_u$ electronic levels. En route back to the $X^1 \\Sigma_g^+$ level, the excited \\htwo\\ may radiate through any of several bound-bound and/or bound-free channels leading to a complex fluorescence emission spectrum.  Located about 20$\\arcmin$ or 1.3 pc from $\\gamma$ Cas, dense gas in the reflection nebula IC 63 is illuminated by a bright UV radiation field.  Line widths of a variety of molecular species are quite narrow \\cite{JVDB94}, suggesting that shock excitation is comparatively unimportant in this environment.  The nebula provides an excellent test case for models of \\htwo\\ fluorescence and other photochemical processes \\cite{JVDBSS95}.  The fluorescence model described in \\citeN{S88} and \\citeN{S89} has been successfully applied in the interpretation of infrared \\htwo\\ fluorescence from the reflection nebulae IC 63 and IC 59 \\cite{LLBJF97}, and to the ultraviolet fluorescence in the band near 1600 \\AA\\ from IC 63 \\cite{S89} and the Taurus molecular cloud \\cite{H94}. In this work we test the model in a previously unexplored band deep in the ultraviolet where the significant majority of the radiated fluorescent power is expected to emerge. ",
        "conclusions": "\\label{conclusions}  We have detected \\htwo\\ fluorescence at wavelengths between 1000 and 1200 \\AA\\ from the bright reflection nebula IC 63 with the Berkeley spectrograph aboard the \\orf-II mission \\cite{Hetal98}. The wavelengths and relative strengths of the fluorescent features within the ORFEUS band agree well with the predictions of the model of \\citeN{S89}. The absolute fluorescent intensity is fainter than an extrapolation from observations at longer ultraviolet wavelengths \\cite{WSBB89} by a factor of ten. The measurements can not be reconciled by differential extinction from a foreground slab of dust (presumably associated with neutral gas) nor by absorption from quiescent \\htwo. Optical depth effects in the fluorescing \\htwo\\ itself, predicted by \\citeN{WSBB89}, remain the most plausible mechanism to explain our observations."
    },
    "9804/astro-ph9804315_arXiv.txt": {
        "abstract": "We present contemporary infrared and optical spectra of the plateau type~II SN~1995V in NGC~1087 covering four epochs, approximately 22 to 84 days after shock breakout. The data show, for the first time, the {\\it infrared} spectroscopic evolution during the plateau phase of a typical type~II event. In the optical region P~Cygni lines of the Balmer series and of metals such as Sc~II, Fe~II, Sr~II, Ca~II and Ba~II lines were identified. The infrared (IR) spectra were largely dominated by the continuum, but P~Cygni Paschen lines and Brackett~$\\gamma$ lines were also clearly seen. The other prominent IR features are confined to wavelengths blueward of 11000~\\AA\\ and include Sr~II 10327, Fe~II 10547, C~I 10695 and He~I 10830~\\AA.  Helium has never before been unambiguously identified in a type~IIp supernova spectrum during the plateau phase.  We demonstrate the presence of He~I 10830~\\AA\\ on days 69 and 85.  The presence of this line at such late times implies re-ionisation.  A likely re-ionising mechanism is $\\gamma$-ray deposition following the radioactive decay of $^{56}$Ni. We examine this mechanism by constructing a spectral model for the He~I 10830~\\AA\\ line based on explosion model s15s7b2f of Weaver \\& Woosley (1993).  We find that this does not generate the observed line owing to the confinement of the $^{56}$Ni to the central zones of the ejecta. In order to reproduce the He~I line, it was necessary to introduce additional upward mixing or ``dredge-up'' of the $^{56}$Ni, with $\\sim$10$^{-5}$ of the total nickel mass reaching above the helium photosphere.  In addition, we argue that the He~I line-formation region is likely to have been in the form of pure helium clumps in the hydrogen envelope. The study of He~I 10830~\\AA\\ emission during the photospheric phase of core-collapse supernovae provides a promising tool for the constraint of initial mixing conditions in explosion models. ",
        "introduction": "Type~IIp (plateau) supernovae form the classic subgroup of the core-collapse supernovae.  They are believed to arise from massive stars (12-25 M$_{\\odot}$) during the red supergiant phase. Early theoretical work by Falk and Arnett (1973) showed that hydrodynamical instabilities should appear in explosions of such massive stars.  As the shock wave propagates through the stellar envelope it sets up density and pressure profiles which can in some cases result in the formation of Rayleigh-Taylor~(RT) instabilities (Chevalier 1976).  In particular, RT instabilities are expected to grow at the interface of the core and the hydrogen envelope because of the large entropy (and density) jump that occurs there (Weaver \\& Woosley 1980).  A direct consequence of these instabilities is that chemical mixing in the ejecta takes place (Bandiera 1984).  Herant \\& Woosley (1994) have studied 2-D simulations of red supergiant explosions over a wide mass range, and found that the growth of hydrodynamic instabilities is highly likely in all cases.  They showed that as the explosion (outgoing) shock plows into the hydrogen envelope, a reverse (ingoing) shock is formed, and between them, RT instabilities grow. Bubbles of hydrogen formed by these instabilities are violently dragged towards the centre of the star by the reverse shock. Simultaneously, compact helium and oxygen clumps advance out into the hydrogen envelope and bubbles of $^{56}$Ni are formed and distributed in the outer parts of the ejecta. It has also been realised that strong dredge-up should result from the neutrino-driven convection close to the neutron star surface (Herant \\& Benz 1992) which has been invoked to account for the conversion of the core-collapse to explosion. This should also produce fast-moving blobs or ``fingers'' of radioactive material which eventually penetrate the outer layers of the supernova.   So far, only SN~1987A has provided us with clear observational evidence for dredge-up in a core-collapse event. This includes the shape of the light curve, the early detection of X-rays and $\\gamma$-rays and the width of the iron lines in the infrared.  However SN~1987A was, of course, only a single event, and a rather unusual one in that it arose from a blue supergiant progenitor. Therefore, we cannot simply assume that similar dredge-up occurs in all other type~II events. Indeed, simulations have shown that differences in progenitor structure can lead to significantly modified hydrodynamical evolution (Herant \\& Benz 1992; Herant \\& Woosley 1994).  Clearly, to establish whether or not deep dredge-up is typical of all core-collapse supernovae, a major step would be to demonstrate dredge-up and mixing in the most-common of all core-collapse events, the type~IIp supernova.  A powerful demonstration of the occurrence of deep dredge-up would be the appearance of radioactive material at the surface at early times. Helium lines arising in the supernova envelope can be used as a tracer of the upwardly-mixed radioactive material. Helium lines are of high excitation. During the first week of the supernova, recombination maintains the populations of the excited levels and so He~I lines are seen.  However after about 10 days, the conditions in the type~IIp atmosphere are such that all the helium will have recombined and de-excited to the ground state.  But, if dredge-up occurs during the explosion, radioactive $^{56}$Ni may reach the outer parts of the supernova envelope at early times.  If it does, the $\\gamma$-rays from its decay ($^{56}$Ni $\\Rightarrow$ $^{56}$Co $\\Rightarrow$ $^{56}$Fe) will excite or re-ionise the helium. Thus the detection of helium lines during the plateau phase (20-120~d post explosion) should imply upward mixing of radioactive material from the core.  Unfortunately no optical He~I lines have ever been unambiguously identified during the plateau phase. However, there are two well-known strong lines in the infrared viz. He~I 10830~\\AA\\ (2s$^{3}$S--2p$^{3}$P$^{0}$) and 20580~\\AA\\ (2s$^{1}$S--2p$^{1}$P$^{0}$) which offer the prospect of testing for dredge-up of radioactive material. This technique was applied by Graham (1988) and Chugai (1991) using the 10830~\\AA\\ line in the SN 1987A at early times. Lucy (1991) invoked upward mixing of $^{56}$Ni to account for strong optical He I lines in the early-time spectra of type Ib supernovae.   To investigate dredge-up in type~IIp supernovae we began in 1995 a programme of infrared and optical spectroscopy of this type of supernova. The data we present here comprise an extensive set of IR/optical spectroscopic observations of the type~IIp SN~1995V, spanning epochs of 22 to 85 days post explosion.  The observations are described in section 2.  In section 3 we compare the observations with a simple spectral synthesis model and discuss the line identifications, especially in the IR. In section 4 we describe the method we used to estimate the amount of dredge-up of $^{56}$Ni. In section 5 we present the results from the comparison of the model with the data, and our estimations for the amount of dredge-up. In section 6 we discuss the implications of this work for our understanding of dredge-up. ",
        "conclusions": "Reproduction by explosion model s15s7b2f of the observed He~I emission is achieved only by invoking substantial additional dredge-up of the $^{56}$Ni from the core.  As expected, the case with no helium clumping requires the greatest dredge-up. However, we argue that the no-clumping case is probably highly unrealistic, since it is difficult to see how one could achieve such a large amount of $^{56}$Ni dredge-up and yet have no pure helium clumps in the same environment.  We deduce, therefore, that there must exist some pure helium clumps in the hydrogen envelope.  As we have shown, with fairly modest clumping of $\\chi_{He}=$0.1 or 0.2, the uniform central core coincides, to within the errors with that of the unmodified s15s7b2f model.  The addition of a steep, power-law density component to this will bring sufficient $^{56}$Ni to the surface to account for the He~I 10830~\\AA\\ emission. Nevertheless, the fraction of the $^{56}$Ni mass which must be dredged up beyond the uniform core is quite substantial. For SN~1995V we conclude that a) a small amount of $^{56}$Ni ($\\sim$10$^{-6}$ M$_{\\odot}$) must have been dredged up to the helium photosphere (v$\\sim$4,250~km/s), and b) clumps of pure helium must have also existed in this region.  High velocities for the decay products of $^{56}$Ni ($\\sim$ 3000 km/s) were also observed in the ejecta of SN 1987A (e.g. Meikle et al. 1993). As shown by Herant \\& Benz (1992), if the $^{56}$Ni is located at the base of the ejecta at t$\\sim$300s it is impossible to accelerate even a small fraction to about 3000 km/s during subsequent instabilities.  In order to achieve such high velocities, it is necessary to invoke outward mixing of the nickel at even earlier times, such as might be caused by neutrino convection. If this occurred, then the later instabilities would carry the nickel to still higher velocities. In order to match the observations of SN1987A, Herant \\& Benz had to premix nickel out to 1.5 M$_{\\odot}$ above the mass cut.  Herant \\& Woosley (1994) studied shock propagation, mixing and clumping in the explosion of red supergiants.  In order to take into account the pre-mixing of $^{56}$Ni during the initiation of the explosion, they diluted the nickel by a factor $\\sim$4 above the mass cut. They then followed the shock propagation and the growth of RT instabilities.  For all progenitors their simulations showed that extensive RT instabilities develop in the ejecta in the wake of the reverse shock from the H/He interface. In contrast to the blue supergiant studies, these instabilities have ample time in which to evolve and completely reshape the ejecta. In spite of this, in all the explosions simulated, nickel did not reach velocities higher than $\\sim$1500 km/s. Similarly, helium did not exceed velocities higher than 2500 km/s. Our results, therefore, indicate that a higher degree of pre-mixing may be required than is invoked in the Herant \\& Woosley models.  Recently, Bazan \\& Arnett (1997) have simulated mixing in core-collapse events. Their simulations include both RT and Richtmeyer-Meshkov instabilities.  These produce much higher velocities for $^{56}$Ni than do R-T instabilities alone. Velocities as high as $\\sim$4000~km/s are predicted. The mass and profile index of the upwardly-mixed $^{56}$Ni derived above could be of considerable value in constraining the initial parameters of these instabilities."
    },
    "9804/astro-ph9804123_arXiv.txt": {
        "abstract": "s{ The X-ray emission from clusters of galaxies is one of the most pursued observational probe to investigate the distribution of dark matter and the related density parameter $\\Omega_0$.  The crucial link to derive the statistics of observables from a dynamical theory is constituted by the physics for the diffuse baryons (or ICP) responsible of the X--ray emission.  Here we present a physical model for the ICP which leads to a definite $L$--$T$ relation. Then we perform a physically based cosmological test, pointing out three cold dark matter universes: a Tilted critical CDM, a flat CDM with $\\Omega_0=0.3$, and an Open CDM with $\\Omega_0=0.5$, which are discussed on the basis of the RDCS survey.  } ",
        "introduction": "Groups and clusters of galaxies constitute cosmic structures sufficiently close to equilibrium and with sufficient density contrast ($\\delta\\approx 2\\, 10^2$ inside the virial radius $R$) as to yield definite observables.  They are dominated by dark matter (hereafter DM), while the baryon fraction is observed to be less than $20$\\%.  The great majority of these baryons are in the form of {\\sl diffuse plasma} (ICP) with densities $n\\sim 10^{-3}$ cm$^{-3}$ and virial temperatures $k\\,T\\sim 5$ keV, and are responsible for powerful X--ray luminosities $L\\sim 10^{44}$ erg/s by optically thin thermal bremsstrahlung.  As the plasma is a good tracer of the potential wells, much better than member galaxies, the X--ray emission is a powerful tool to investigate the mass distribution out to moderate and high redshifts.  The ICP temperature directly probes the height of the potential well, with the baryons in the  role of  mere tracers; on the other hand, the luminosity with its strong dependence on density ($L\\propto n^2$) reliably probes the baryonic content and distribution. Statistically, an average $L$--$T$ correlation is observed along with substantial scatter, and this provides the crucial link to relate the X-ray luminosity functions with the underlying statistics of the DM.  A physical model for the diffuse baryons is difficult to achieve. In fact, the simple self similar model (Kaiser 1986), which assumes the ICP amount to be proportional to the DM's at all $z$ and $M$, leads to a relation $L\\propto T^2$, conflicting with the observed correlation for rich clusters. The latter is close to $L\\propto T^{3.5}$ (David et al. 1993; Mushotzky \\& Scharf 1997).  Here we propose a physical model for baryons, which leads to a prediction for the $L$--$T$ relation (see Cavaliere, Menci \\& Tozzi 1997, CMT97) and allows a non parametrical approach to the search for cosmological parameters.  The results, presented in \\S 3, are a synthesis from Cavaliere, Menci \\& Tozzi (1998, CMT98). ",
        "conclusions": "We presented a physically based approach to cosmological tests with clusters of galaxies.  We describe the X--ray emission from clusters with a specific model for the diffuse baryons.  In this sense this approach is alternative to the parametrical approach (see Borgani this meeting).  The results of our model depend on two parameters, the external temperature $T_1$ and density $n_1$, which are {\\sl not} free. Specifically, we use for $T_1$ the range $0.1\\div 0.8 $ keV provided by the stellar preheating.  The value of $n_1$ for rich clusters is related to the DM density by the universal baryonic fraction.  Thus we compute the expression of the bolometric luminosity for a given temperature.  The average of the square of the density jump factor $\\langle g^2 \\rangle$ over the merging histories coupled with $\\beta(T)$ is what gives to the statistical $L-T$ correlation the curved shape shown in fig. \\ref{lt}b. In addition, our approach predicts an intrinsic {\\it variance} of dynamical origin due to the different merging histories, and built in the factor $g^2$.  With the ICP state so described, we proceeded to constrain the cosmological parameters. After the observations by Rosati et al. (1998), we have computed the X-ray observables for groups and clusters of galaxies.  On the basis of local data, the set of acceptable CDM universes is restricted to three disjoint domains: $\\Omega=1$ for the Tilted CDM with high baryon content; $\\Omega_o\\simeq 0.5$ for standard CDM; $\\Omega_o\\approx 0.3$ for CDM in flat geometry.  However, only the TCDM and the $\\Lambda$CDM universes give acceptable faint counts. As an overall remark, a common feature of all the above universes is constituted by some excess in the counts.  This may indicate some non--trivial incompleteness in the canonical hierarchical clustering, worth keeping under scrutiny.  We recall that in the adiabatic models for the ICP (Evrard \\& Henry 1991, Kaiser 1991) the evolution of the $L-T$ relation is reduced or even negative, thus alleviating the excess.  However, the anti--evolution required in OCDM would be very difficult to justify (the adiabatic models are largely discussed in CMT98).  Now the question is: to what extent enlarging the data base on X--ray clusters will help in further constraining cosmology?  We argue that the variance intrinsic to the hierarchical clustering, and amplified by the ICP emissivity, sets an effective limitation.  Richer, faint surveys will hardly provide a sharper insight into cosmology unless one reduces both the uncertainty concerning $\\sigma_8$ and the larger one concerning $L_{o}$. However, we stress that such efforts will find soon a more proper aim than constraining $\\Omega_o$.  This is because MAP, and subsequently PLANCK, will accurately measure on still linear scales not only the perturbation power spectrum but also directly $\\Omega_o$.  Once the cosmological framework has been fixed, the study of groups and clusters in X-rays will resume its proper course, that is, the physics of systems of intermediate complexity which is comprised of the DM and of the ICP component."
    },
    "9804/astro-ph9804065_arXiv.txt": {
        "abstract": "We present an empirical method which measures the distance to a Type Ia supernova (SN~Ia) with a precision of $\\sim$ 10\\% from a single night's data.  This method measures the supernova's age and luminosity/light-curve parameter from a spectrum, and the extinction and distance from an apparent magnitude and color.  We are able to verify the precision of this method from error propagation calculations, Monte Carlo simulations of well-sampled SNe~Ia, and the Hubble diagram of scarcely observed supernovae.  With the reduction in telescope time needed, this method is three to four times more efficient for measuring cosmological parameters than conventional light-curve based distance estimates. ",
        "introduction": "The explosion of a Type Ia supernova (SN~Ia) is a catastrophic phenomenon veiled in layers of complexity.  Recent efforts to monitor these events have led to an increased ability to predict, if not fully understand, the stages of SN Ia evolution.  The model for the photometric history of SNe~Ia has been refined from a homogeneous description \\cite[]{brunophd,branch_miller,ts95} to one which characterizes a relation between peak luminosity and light-curve shape \\cite[]{philm15,hametal95,hametal96,rpk95,kim_stretch98}. The slower, broader light-curves are intrinsically brighter at peak than the faster, narrower light-curves. Recognizing and exploiting such relations has led to a renaissance in the use of SNe~Ia as extragalactic distance indicators.  Extending luminosity/light-curve relations to multiple passbands separates the competing effects of dust, intrinsic differences, and distance on the light of SNe~Ia \\cite[]{rpk96}.  Distances with 5-10\\% uncertainty can be obtained using the light-curve shapes of well-observed supernovae.  The optical spectra of SNe~Ia are rich in information [see \\cite{fil97} for a review].  Many of the elements synthesized and ejected in the explosion have been identified despite the blending of their high-velocity profiles \\cite[]{bran81b,nugphd}.  In addition, the relative strengths of some spectral features have been shown to correlate with SN~Ia peak luminosity \\cite[]{nugseq95}.  As the supernova evolves, predictable casts of features appear and disappear, illuminated by the photosphere's recession through the synthesized layers.  The temporal evolution of these features is sufficiently reliable to be used as a clock to determine the current age of a SN~Ia to a precision of 1-2 days \\cite[]{mink39,riess_age97}.  Unfortunately, supernovae occur without warning, making it difficult to collect the observations necessary to measure their distances. Observing an unscheduled event in up to four filters many times over the course of $\\sim$ 100 days is a time consuming and logistically formidable task. The observing record for a typical SN~Ia is quite fragmentary.  Following this process, it will be many years of work to gather the number of SN~Ia distances necessary to put strong limits on cosmological parameters. Even at high redshifts ($z \\geq 0.3$), where a strategy for batch detections of multiple supernovae has made it possible to schedule supernova discoveries and their follow up \\cite[]{perl97}, difficulties arise. Since these observations require the largest telescopes, the light-curves are typically more poorly sampled than the nearby ones caught at a similar phase.  Yet, from a single night's observations, a SN~Ia's spectrum and photometric magnitudes can reveal its age, intrinsic luminosity, extinction, and apparent brightness.  From this information one can estimate the distance to a supernova without further observations (except for the possible need of a galaxy image to subtract the host's light).  Here we explore this possibility with two independent sets of SNe~Ia. We describe this technique in \\S 2 and its expected uncertainty in \\S 3.  In \\S 4 we apply it to randomly selected snapshots of extensively observed SNe~Ia to determine the precision of such distance estimates. In \\S 5 we construct the Hubble diagram of ``cast-off'' SNe~Ia: objects which were observed only once or a few times.  We extend the application of this method in \\S 6 to SNe~Ia with $0.2 \\leq z \\leq 0.83$.  In \\S 7 we discuss variations of this technique and its leverage on estimating cosmological parameters. ",
        "conclusions": "In principle, enough information can be garnered from a single supernova spectrum and photometric epoch to estimate the distance to a SN~Ia.  In practice, the results of \\S 3, 4, and 5 suggest this method produces distances having a precision of $\\sim$ 10\\%, with variations that are a function of the quality of the data and the age of the supernova.  Depending on the amount of host galaxy contamination, it may be necessary to obtain spectra and images of the host galaxy after the supernova has faded.  The snapshot distance method employs the same luminosity and extinction corrections used in the MLCS method of Riess, Press, and Kirshner 1996a (and more recently updated in Riess et al. 1998b). We find no significant offset between the distance estimates of the two methods and only a modest reduction in precision for the snapshot distance method.  There are three more limited versions of a distance method using single epoch SN Ia observations which reveal the utility of luminosity and extinction corrections (see Table~\\ref{discomp_tab} and Figure~\\ref{show}).  These variants employ the SFA measurement but lack the luminosity correction, the extinction correction, or both.  Disregarding individual luminosities and light-curve shapes predicted by the spectral ratios, we fit homogeneous, fiducial templates to the photometric epoch.  Further, we discarded our estimate of the extinction from the color excess.  We measured the resulting ``standard candle'' distances to SNe~Ia using the Monte Carlo technique described in \\S 4.  As seen in Figure~\\ref{show}, the distribution of dispersions has a mean of 0.35 mag, a value consistent with previous SNe~Ia distance estimates which assume SN Ia light curve homogeneity and do not correct for extinction \\cite[]{st93,ts95,hametal95,hametal96,rpk95,rpk96,branch_miller,vauetal95}. These distances are also 15\\% {\\it greater} in the mean (or smaller in the implied Hubble constant) than either MLCS or snapshot distances, consistent with other comparisons of distance estimates which assume homogeneity instead of heterogeneity of SNe~Ia \\cite[]{ts95,hametal95,hametal96,rpk95,rpk96}.  To simulate the effect of single filter information, we used our luminosity correction without an extinction correction.  This procedure results in a mean dispersion of 0.25 mag and distances which are only 3\\% greater than those obtained from MLCS.  This result, though better than the standard candle method, is still worse than the complete snapshot procedure.  A final variant is to disregard a luminosity/light-curve shape correction but maintain an extinction correction from the color excess.  Such a method using light-curves was proposed by \\cite{vdb95} to account for both intrinsic luminosity differences as well as absorption by dust.  This method takes advantage of the coincidental near-agreement between the standard reddening law and the relation between intrinsic color and luminosity to correct for both extinction and luminosity differences.  \\cite{rpk96} have noted that while both sources of luminosity variation affect the SN color in the same direction, the specific ratios of the luminosity difference to color difference are not precisely the same for extinction and intrinsic SN Ia variation.  Monte Carlo simulations of this method combined with a SFA measurement give a mean dispersion of 0.19 mag and distances which are 5\\% greater in the mean than those of the MLCS method.  Despite the low dispersion of this method, we are suspicious of the distances it predicts.  The distribution of dispersions obtained from our Monte Carlo simulation (see Figure~\\ref{show} and Table~\\ref{discomp_tab}) is more skewed than any other, including an asymmetric ``tail'' encompassing dispersions greater than 0.3 mag.  We believe that for SNe~Ia with only moderate amounts of extinction or whose luminosities are similar to those of typical SNe~Ia, this method has merit.  Yet for very red SNe~Ia, this method can predict distances which are systematically and considerably in error due to the inability to distinguish between absorption by dust and intrinsic variation.  The snapshot method predicts distances which agree in the mean with only a moderate reduction in precision from light-curve shape methods. Yet because of the greatly diminished expense in data collection, this method can be more effective for problems which benefit equally well from a high {\\it quantity} of SNe~Ia distances as from the {\\it quality} of those distances.  Two such applications are mapping the nearby peculiar velocity field and determining the cosmological parameters which dictate global geometry.  Recent attempts to map the cosmic velocity field with SNe~Ia \\cite[]{riess_beta97} suffer from dilute spatial sampling.  Replete peculiar velocity maps could reveal the influence of matter fluctuations and constrain the matter content of the local Universe. Nearby, many more SNe~Ia are discovered than can be regularly monitored.  By decreasing the observational requirements of each SN Ia, it should be possible to increase the sampling of the local velocity field.  The light-curves employed by \\cite{hametal96} and \\cite{rpk96} are typically sampled for 10 to 15 epochs.  The observational demands increase as the supernova rapidly fades.  With the telescope time invested in a single set of SN Ia multi-color light-curves, sufficient data for 15 to 25 SNe~Ia snapshot distances could be gathered.  Accounting for the inherent distance uncertainties, telescope time spent collecting snapshot distance data is 3 to 4 times more efficient than time spent collecting light-curves for distance estimates.  Efforts to measure the cosmological parameters $\\Omega_M$ and $\\Omega_\\Lambda$ from distant SNe~Ia could also profit from snapshot distances.  Systematic searches for SNe~Ia at z $\\geq 0.3$ have yielded a plethora of objects \\cite[]{perliau95a,schetal95}. Combining a new generation of large telescopes with $\\sim$ 1 degree fields of view and multi-fiber spectrometers with the snapshot method could allow SNe~Ia distances to be gathered in batch at an unprecedented rate.  At the current rate of discovery, a night spent searching five 1-degree fields followed by a night collecting spectra of the candidates with a multi-fiber spectrometer could yield $\\sim$ 50 SNe~Ia distances \\cite[]{schmidt97,rate_96}. Repeating this process every new moon could yield up to $\\sim$ 600 distances a year.  At this rate of accumulation, it should be possible to convincingly separate the effects of various sources of energy density on the redshift-magnitude relation \\cite[]{omol_95}.  A more optimal method for measuring SN Ia distances would employ both the predictive power of SN Ia light and color curve shapes with that of SN Ia spectra.  Such a method would replace the distinction between a snapshot distance and a light-curve distance with a distance estimate which makes the most economical use of all available SN Ia observations.  Using the tools described in Riess, Press, and Kirshner 1996a, Nugent et al. 1995, Riess et al. 1997, Riess et al. 1998b, and this paper, such a method appears to be quite feasible.   \\bigskip  This work was supported by the NSF through grant AST--9417213 to A.V.F. and AST-9528899 and AST-9617058 to R.P.K., by the Miller Institute for Basic Research in Science through a fellowship to A.G.R., and by the Director, Office of Computational and Technology Research, Division of Mathematical, Information, and Computational Sciences of the U.S. DoE under contract number 76SF00098 to P.E.N. Some of the calculations presented in this paper were performed at the National Energy Research Supercomputer Center (NERSC), supported by the U.S. DoE.  We thank Stephan Benetti and George Djorgovski for allowing us to use their spectra of SNe Ia prior to publication, and Bruno Leibundgut for suggestions that helped improve this paper.  \\appendix"
    },
    "9804/hep-ph9804285_arXiv.txt": {
        "abstract": "The expected proton and neutrino fluxes from decays of massive metastable relic particles are calculated using the HERWIG QCD event generator. The predicted proton spectrum can account for the observed flux of extremely high energy cosmic rays beyond the Greisen-Zatsepin-Kuzmin cutoff, for a decaying particle mass of ${\\cal O}(10^{12})$~GeV. The lifetime required is of ${\\cal O}(10^{20})$~yr if such particles constitute all of the dark matter (with a proportionally shorter lifetime for a smaller contribution). Such values are plausible if the metastable particles are hadron-like bound states from the hidden sector of supersymmetry breaking which decay through non-renormalizable interactions. The expected ratio of the proton to neutrino flux is given as a diagonistic of the decaying particle model for the forthcoming Pierre Auger Project. ",
        "introduction": "It has been known for some time that interactions on the 2.73 K blackbody cosmic microwave background (CMB) will severely degrade the energies of cosmic ray nucleons with energies beyond $\\sim5\\times10^{19}\\ev$ --- the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{gzk}. It was therefore very surprising when the Fly's Eye atmospheric fluorescence detector reported the observation of an extremely high energy cosmic ray (EHECR) event with an energy of $(3.0\\pm0.9)\\times10^{20}\\ev$ \\cite{flyseye}. This was followed by the detection of a $(1.7-2.6)\\times10^{20}\\ev$ event by the AGASA air shower array \\cite{agasa}. These discoveries substantiated earlier claims from the Volcano Ranch \\cite{vr}, Haverah Park \\cite{hp} and Yakutsk \\cite{yak} air shower arrays that cosmic rays do exist beyond the GZK cutoff. About a dozen such events are now known. Detailed accounts of the data may be found in recent reviews \\cite{reviews}.  In Figure~\\ref{fig1} we show the EHECR spectrum for energies exceeding $10^{18}\\ev$ \\cite{spec}; note that the fluxes have been multiplied by $E^3$. It is believed that cosmic rays with energies up to $\\sim5\\times10^{18}\\ev$, the so-called `ankle', are predominantly of galactic origin, possibly accelerated by the Fermi mechanism in supernova remnants \\cite{books}. Above this energy, the spectrum flattens and the composition changes from being mostly heavy nuclei to mostly protons. Such a correlated change in the spectrum and composition was first established by the Fly's Eye experiment \\cite{flyseye} and Figure~\\ref{fig1} shows their suggested two-component fit to the data. The new component which dominates at energies beyond $\\sim5\\times10^{18}\\ev$ is isotropic and therefore cannot possibly originate in the galactic disk \\cite{agasa2,lc95}. However it also extends well beyond the GZK cutoff raising serious problems for hypothetical extragalactic sources. Because of the rapid energy degradation at these energies through photo-pion production on the CMB, such sources must exist within $\\sim500\\mpc$, in fact within $\\sim50\\mpc$ for the highest energy Fly's Eye event \\cite{cronin}. For heavy nuclei, the energy loss is less severe according to a revised calculation \\cite{heavy} so the range may extend upto $\\sim100\\mpc$.  General arguments \\cite{greisen,hillas} provide correlated constraints on the magnetic field strength and spatial extent of the region necessary to accelerate particles to such high energies and these requirements are barely met by likely astrophysical sites such as active galactic nuclei and the `hot spots' of radio galaxies \\cite{cr}. Moreover there are few such sources close to us and no definite correlations have been found between their locations and the arrival directions of the most energetic events \\cite{source,agasa2}. It has been speculated that gamma-ray bursts which too are isotropically distributed, may be responsible for EHECRs \\cite{grb}. However since these are at cosmological distances, one would expect to see the GZK cutoff in the cosmic ray spectrum contrary to observations (cf. ref.\\cite{grbtest}).  Some of the above arguments may be evaded if the EHECR events are due not to nucleons but neutral particles such as photons and neutrinos. Although high energy photons also suffer energy losses in traversing the CMB and the extragalactic radio background, there is no threshold effect which would cause a cutoff near the GZK value \\cite{photon}. However the observed shower profile of the highest energy Fly's Eye event \\cite{flyseye} argues against the primary being a photon since it would have interacted on the geomagnetic field and started cascading well before entering the atmosphere \\cite{notphoton}. The observed events are also unlikely to be initiated by neutrinos as they all have incident angles of less than $40^\\circ$ from the zenith and thus too small a path length in the atmosphere for interactions \\cite{gqrs96}. This argument may be evaded if neutrinos become strongly interacting at high energies due to new physics beyond the Standard Model \\cite{neutrino,hongmo}, but such proposals are found not to be phenomenologically viable \\cite{bhg98} (although this is disputed \\cite{nu}). (Alternatively, the propagating high energy neutrinos could annihilate on the relic cosmic neutrino background, assumed to have a small mass of ${\\cal O}(0.1)$~eV, to make hadronic jets within the GZK zone \\cite{weiler}.) Other exotic possibilities have been suggested, e.g. monopoles \\cite{wk96}, stable supersymmetric hadrons \\cite{farrar} and loops of superconducting cosmic string (`vortons') \\cite{bp97}. However these possibilities have many phenomenological problems \\cite{mn98,ber} and we do not discuss them further.  Thus one is encouraged to seek `top-down' explanations for EHECRs in which they originate from the decay of massive particles, rather than being accelerated up from low energies. The most discussed models in this connection are based on the annihilation or collapse of topological defects such as cosmic strings or monopoles formed in the early universe \\cite{hill,td,bs95,bv97}. When topological defects are destroyed their energy is released as massive gauge and Higgs bosons which are expected to have masses of ${\\cal O}(10^{16})\\gev$ if such defects have formed at a GUT-symmetry breaking phase transition. The decays of such particles can generate cascades of high energy nucleons, $\\gamma$-rays and neutrinos. A more recent suggestion is that EHECRs arise from the decays of metastable particles with masses $m_X\\sim10^{13}-10^{16}\\gev$ which constitute a fraction of the dark matter \\cite{bkv97}. These authors suggest that such particles can be produced during reheating following inflation or through the decay of hybrid topological defects such as monopoles connected by strings, or walls bounded by strings. The required metastability of the particle is ensured by an unspecified discrete symmetry which is violated by quantum gravity (wormhole) effects. Another suggestion is that the long lifetime is due to non-perturbative instanton effects \\cite{kr97}. In ref.\\cite{fkn97}, a candidate metastable particle is identified in a $SU(15)$ GUT.  A generic feature of these `top-down' models is that the EHECR spectrum resulting from the decay cascade is essentially determined by particle physics considerations. Of course the subsequent propagation effects have astrophysical uncertainties but since the decays must occur relatively locally in order to evade the GZK cutoff \\cite{bkv97}, they are relatively unimportant. Thus although the proposal is speculative, it is possible, in principle, to make reliable calculations to confront with data. In this work we consider another possible candidate for a relic metastable massive particle \\cite{ben98} whose decays can give rise to the observed highest energy cosmic rays. First we discuss (\\S~\\ref{crypton}) why this candidate, which arises from the hidden sector of supersymmetry breaking, is perhaps physically better motivated than the other suggested relics. We then undertake (\\S~\\ref{decay}) a detailed calculation of the decay cascade using a Monte Carlo event generator to simulate non-perturbative QCD effects. This allows us to obtain a more reliable estimate of the cosmic ray spectrum than has been possible in earlier work on both topological defect models \\cite{td} and a decaying particle model \\cite{bkv97}. We confront our results with observations and identify the mass and abundance/lifetime required to fit the data. We conclude (\\S~\\ref{concl}) with a summary of experimental tests of the decaying particle hypothesis. ",
        "conclusions": "We have investigated the hypothesis that the highest energy cosmic rays, in particular those observed beyond the GZK cutoff, arise from the decay of massive metastable relic particles which constitute a fraction of the dark matter in the galactic halo. To simplify computations (using the HERWIG Monte Carlo event generator) we have considered only decays into $q\\bar{q}$ pairs with unit branching ratio. Comparison with experimental data indicates that a decaying particle mass of ${\\cal O}(10^{12})\\gev$ is required to fit the spectral shape while the absolute flux requires a lifetime of ${\\cal O}(10^{20})\\yr$ if such particles contribute the critical density. The predicted decay spectra may be somewhat altered if 3-body decays and other final states (e.g. supersymmetric particles \\cite{bk97}) are considered. However our conclusions regarding the preferred mass and relic abundance/lifetime of the decaying particle are unlikely to be affected. In particular it would appear that the approximations used to calculate the particle spectra in previous studies of decaying topological defects \\cite{td} and hypothetical massive particles \\cite{bkv97} were not sufficiently accurate. Our work indicates that the topological defect model is disfavoured unless the mass of the decaying gauge bosons is less than about $10^{13}\\gev$, which is well below the unification scale of $\\sim10^{16}\\gev$. (A similar conclusion is arrived at by independent arguments in refs.\\cite{bss97,vah98}.) By contrast, cryptons from the hidden sector of supersymmetry breaking have a mass of the required order, as well as a decay lifetime which is naturally suppressed. However their relic abundance is difficult to estimate reliably, although we have argued that it may be cosmologically interesting.  The primary intention of this work is to attempt to quantify the decaying particle hypothesis in a manner which is of interest to experimentalists. We have therefore computed the expected neutrino to proton ratio as a function of energy since this is an important test of competing hypotheses for forthcoming experiments, in particular the Pierre Auger project \\cite{auger}. Of course our cleanest prediction is that the cosmic ray spectrum should cut off just below the mass of the decaying crypton, at $\\sim3\\times10^{20}\\ev$. Moreover, with sufficient event statistics it should be possible to identify the small anisotropy which should result from the distribution of the decaying particles in the Galactic halo \\cite{dt98}. Thus although the hypothesis investigated here is very speculative, it is nevertheless testable. Perhaps Nature has indeed been kind to us and provided a spectacular cosmic signature of physics well beyond the Standard Model."
    },
    "9804/astro-ph9804175_arXiv.txt": {
        "abstract": "s{We are contructing an interferometric telescope, the Very Small Array, to study the cosmic microwave background on angular scales 0.2--$4.5^{\\circ}$. The physical layout and electronic design of the telescope are optimised to give maximum protection from systematic effects, while still providing sufficient sensitivity to make high signal-to-noise images. A prototype single baseline is currently being tested, with scientific results expected during 2000.} ",
        "introduction": "It is widely accepted that the majority of the cosmological results that might be obtained from the cosmic microwave background (CMB) will come from accurate measurements of the CMB power spectrum over the region of the accoustic peaks, that is in the range $100 < l < 2000$ (where $l$ is the spherical harmonic multipole). We have therefore been constructing an instrument, the Very Small Array (VSA), to measure the power spectrum (and provide images) in precisely this angular range. We have argued previously\\cite{jones} that interferometers are very well suited for ground-based CMB measurements, given their relative immunity to the atmosphere and other systematic effects compared with switched-beam experiments. Here we will describe some of the more detailed design considerations of the array, and review the progress of the project so far. ",
        "conclusions": ""
    },
    "9804/astro-ph9804205_arXiv.txt": {
        "abstract": "We present HST WFPC2 observations in three bands (F450W=B, F467M and F814W=I) of a group of three galaxies at $z=2.8$ discovered in a ground-based narrow-band search for \\lya emission near the $z=2.8$ quasar \\pks. One of the galaxies is a damped \\lya (DLA) absorber and these observations bear on the relation between the DLA clouds and the Lyman-break galaxies and the stage in the evolution of galaxies they represent. We describe a procedure for combining the undersampled WFPC2 images pointed on a sub-pixel grid, which largely recovers the full sampling of the WFPC2 point spread function (psf). These three galaxies have similar properties to the Lyman-break galaxies except that they have strong \\lya emission. The three galaxies are detected in all three bands, with average $m_B\\sim26$, $m_I\\sim25$. Two of the galaxies are compact with intrinsic (i.e. after correcting for the effect of the psf) half-light radii of $\\sim0.1$ arcsec ($0.4 h^{-1}$ kpc, $q_{\\circ}=0.5$). The third galaxy comprises two similarly compact components separated by 0.3 arcsec. The HST images and a new ground-based \\lya image of the field provide evidence that the three galaxies are more extended in the light of \\lya than in the continuum. Combined with the evidence from the \\lya line widths, previously measured, this suggests that we are measuring the size of the surface of last scattering of the escaping resonantly-scattered \\lya photons. The measured impact parameters for this DLA galaxy (1.17 arcsec), for a second confirmed system, and for several candidates, provide a preliminary estimate of the cross-section-weighted mean radius of the DLA gas clouds at $z\\sim 3$ of $<13 h^{-1}$ kpc, for $q_{\\circ}=0.5$. The true value is likely substantially smaller than this limit as DLA clouds at small impact parameter are harder to detect. Given the observed sky covering factor of the absorbers this implies that for $q_{\\circ}=0.5$ the space density of DLA clouds at these redshifts is more than five times the space density of spiral galaxies locally, with the actual ratio probably considerably greater. For $q_{\\circ}=0.0$ there is no evidence as yet that DLA clouds are more common than spiral galaxies locally. We summarise evidence that filamentary structures occur in the distribution of galaxies at high redshift. ",
        "introduction": "By studying typical galaxies at high redshift, here $z>2$, we can record the origins of normal galaxies such as our own Milky Way. Schmidt \\cite{sc65} was the first to observe a high$-$redshift galaxy when he obtained a spectrum of 3C9, a radio$-$loud quasar of redshift $z=2.01$. Normal star$-$forming high$-$redshift galaxies are about 1000 times fainter, and remained undiscovered until the present decade. A small number of candidates (some of which may be active galaxies) have been identified with 4$-$metre telescopes, through the detection of Ly$\\alpha$ emission (e.g. Steidel, Sargent, and Dickinson 1991, Lowenthal et al. 1991, M\\o ller and Warren 1993, Pascarelle et al. 1996, Francis et al. 1995). Another approach has been to employ deep broad$-$band imaging to identify candidates by the expected Lyman$-$limit discontinuity in their spectra (Steidel and Hamilton 1992). The brightest Lyman$-$break galaxies have $m_R\\sim 24$ and spectroscopic confirmation has only become feasible with the completion of the Keck 10$-$metre telescope. The report by Steidel et al. (1996) on the results of spectroscopic observations of Lyman$-$break candidates really marked the beginning of the statistical study of high-redshift galaxies. They were able to confirm redshifts for 15 star$-$forming galaxies in the range $3.0<z<3.5$. They discovered that Lyman$-$break galaxies generally show weak Ly$\\alpha$ emission, but that the rest$-$frame ultra$-$violet spectra may be recognised by the presence of absorption lines characteristic of the spectra of nearby star$-$forming galaxies. Madau et al. (1996) applied the same methodology to the Hubble Deep Field data to extend the results to fainter magnitudes, and their analysis currently provides the best lower limits to the integrated star formation rate in galaxies at $z>2$.  Complementary to the searches for starlight from high$-$redshift galaxies have been the analyses of absorption lines in the spectra of high$-$redshift quasars. These studies have provided measurements of the mass density of neutral hydrogen in the universe (e.g. Wolfe et al. 1986, Lanzetta et al. 1991), and the metallicity of the gas (e.g. Pettini et al 1994), and how these quantities have changed with redshift. The analysis of Pei and Fall (1995) of the absorption--line data reconstructs the global history of star formation, gas consumption, and chemical enrichment, accounting in a self--consistent way for the effects of the progressive extinction of the background quasars due to dust as star formation proceeds.  The advantage of the absorption--line approach to the history of star formation is that it is global in nature as all the neutral gas at any redshift is directly observed. With the deep imaging studies it is necessary to extrapolate below the survey flux limit to measure the total star formation rate.  The weakness of the global approach however is that it tells us nothing about the sizes and morphologies of the galaxies in which the star formation occurs. Deep imaging of DLA absorbers provides the missing information. In other words if we imaged the absorbers we could combine the data of faint galaxy surveys (luminosities, sizes, shapes) with the information on the gas obtained from the spectroscopic studies (column densities, chemistry), allowing a more detailed comparison with theories of galaxy formation.  Unlike for MgII absorbers (e.g. Bergeron and Boiss\\'e 1991, Steidel, Dickinson, and Persson 1994a) little progress has been made in such a programme of imaging of DLA systems. There are only two published unambiguous detections, discussed below, as well as a small number of candidate counterparts of DLA absorbers (at both low redshift $-$ Steidel et al. 1994b, Le Brun et al. 1997 $-$ and high-redshift $-$ Arag\\'{o}n-Salamanca et al. 1996). In this paper we report on the results of 30 orbits of imaging observations with the {\\it Hubble Space Telescope} (HST) of a group of three galaxies at $z=2.81$, named S1, S2, S3, one of which (S1) is a damped \\lya absorber. The three galaxies lie in the field of the quasar \\pks, and were detected in the light of \\lya emission by M\\o ller and Warren (1993, hereafter Paper I). A detailed discussion of their nature was presented by Warren and M\\o ller (1996, hereafter Paper II \\footnote{We note here that recent spectroscopic observations by Ge et al. (1997) confirm our conclusion that the \\lya emission from the DLA absorber is due to star formation rather than photoionisation by the quasar.}). The HST images provide measures of their sizes, magnitudes, and colours.  The other high$-$redshift damped system that has been successfully imaged is the absorber at $z=3.150$, towards the quasar $2231+131$, observed by Djorgovski et al. \\cite{dj96}. In this paper we intercompare the measured properties of these two absorbers, the two companion galaxies S2 and S3, and the population of Lyman$-$break galaxies, to draw conclusions about the space density of DLA absorbers, their structure, and the relation between the DLA absorbers and the Lyman-break galaxies.  The layout of the rest of the paper is as follows. The HST observations of the field towards \\pks are described in Section 2, as well as new ground$-$based narrow$-$band observations of the field. The three galaxies are very small, and the HST images were dithered with a half-integer pixel step in order to improve the image sampling. The algorithm for combining the images is outlined in this section. In Section 3 we present the reduced HST images. In Section 4 we provide the results of aperture photometry and profile fitting of the HST and ground-based images. Finally, in Section 5 we discuss these results and their implications for our understanding of the nature of DLA absorbers and Lyman-break galaxies. ",
        "conclusions": "\\subsection{Continuum emission}  The three galaxies S1, S2, S3 are similar in their properties to members of the recently detected population of Lyman-break galaxies, except that they have strong \\lya emission, with mean restframe ${\\rm EW=60\\AA}$. For example the B and I magnitudes for our three sources $m_B\\sim 26$, $m_I\\sim 25$, are within the range of magnitudes of the Lyman-break galaxies of redshift $z\\sim3$ observed by Steidel et al. (1996) and Lowenthal et al. (1997). The average half-light radius of the four objects listed in Table 4 is 0.1 arcsec. This is smaller than the value quoted for the Lyman-break galaxies (Giavalisco et al. 1996, Lowenthal et al. 1997), by a factor of two to three. However we have broken S3 into two sub components and measured the radii for each subcomponent. In addition the radii quoted in Table 4 have been corrected for the effect of the psf.  Young star-forming galaxies should have approximately constant flux per unit frequency over the restframe ultraviolet and optical regions of the spectrum. The expected colour of a source at $z=2.8$ with a power-law continuum varying as $f_{\\nu}\\propto\\nu^0$, after accounting for typical absorption in the \\lya forest, is $m_B-m_I=0.66$. As seen in Table 3 the $m_B-m_I$ colours for the three sources S1, S2, S3 are consistent with this value, for both the small and large apertures.  The fact that S1 has properties similar to the Lyman-break galaxies suggests that DLA absorbers and Lyman-break galaxies may commonly be associated with each other. Lowenthal et al (1997) have also made a connection between the two populations by suggesting that the broad \\lya absorption line seen in several cases in the spectra of Lyman-break galaxies is a damped line from neutral hydrogen in the galaxy. However the absorption lines seen in their average spectrum and in the relatively high S/N spectrum presented by Ebbels et al. (1996) are not optically thick in the line centre. The absorption--line profiles are difficult to interpret because the background source is extended.  It is even possible that in some cases the absorption line is stellar absorption in the spectra of late B stars and that the galaxies are observed in a post-burst phase, rather than actively forming stars (see Valls-Gabaud 1993) i.e. for some cases this may be the explanation for the weakness of the \\lya emission line in the Lyman-break galaxies, as opposed to dust.  High-redshift galaxies with strong \\lya emission, like the three sources near \\pks, may therefore be younger than many of the Lyman-break galaxies.  \\subsection{Extended \\lya emission}  In this subsection we summarise the evidence that the \\lya emission from the three sources is more extended than the continuum emission.  The evidence for S2 and S3 comes from an examination of the $m_B-m_M$ and $m_B-m_N$ colours, summarised in Table 3. The bottom row in Table 3 provides the predicted $m_B-m_M$ colours for each source computed, as detailed in Section 4.1.2, from the B and N total magnitudes. The small aperture $m_B-m_M$ colours are smaller than the predicted values, as would be the case if the \\lya emission is more extended than the continuum. For S2 the difference is significant at $2.8\\sigma$, and for S3 at $2.3\\sigma$. Notice also that the large aperture $m_B-m_M$ colours are larger than the small aperture colours, which is consistent with the hypothesis of extended \\lya emission. A broad absorption line in the M passband could only explain part of the discrepancy between the measured and predicted colours. For example for S2, simply removing $50{\\rm\\AA}$ of the continuum flux within the M band changes the $m_B-m_M$ colour by only 0.1 mag.  The total B magnitude for S1 (i.e. the large-aperture measurement from the HST image, Table 4) is not reliable because of the uncertainty associated with the psf subtraction. This prevents us from usefully making the same comparison between predicted and measured $m_B-m_M$ colours for S1. However there is an indication that S1 may also be more extended in \\lya. The measurements of the half-light radius of S1 from the ground-based \\lya image yielded values of $0.51^{+0.08}_{-0.08}$ arcsec for the exponential profile and $0.66^{+0.25}_{-0.19}$ arcsec for the de Vaucouleurs profile (Section 4.2.2). These values are significantly larger than the half light radius of the region of continuum emission of $0.13\\pm0.06$ arcsec (Table 4), measured from the HST image. Such a diffuse contribution of \\lya to the M-band flux would be hard to detect, and could have been partially subtracted in attempting to remove large-scale residuals from the psf subtraction (Section 4.2.1). However we cannot rule out the possibility that there is also a similar diffuse contribution to the continuum flux.  The argument for extended \\lya is made visually in Figure 4. Here we compare the M-band images after subtraction of the computed contribution of the continuum (i.e. showing the \\lya emission only), against how the objects would have appeared if the \\lya light profile were the same as the continuum profile.  \\begin{figure} \\vspace{8.5cm} \\special{psfile=Reffig_1.ps angle=0 voffset=-52 hoffset=-14 vscale=40 hscale=40} \\caption[ ]{Illustration of the evidence for extended \\lya emission for, from top to bottom, S1, S2, S3. Each panel shows the same 2 arcsec by 2 arcsec region as Figure 3, here smoothed by a 0.15 by 0.15 arcsec boxcar filter to enhance faint features. The left column shows the summed image from all three filters. The middle column shows the M-band images with the computed contribution of the continuum subtracted. This image, then, contains only the \\lya line component. In the right column we show how the \\lya-only image would have appeared if the \\lya profile were identical to the continuum profile.} \\end{figure}  To summarise, there is evidence that the regions of \\lya emission from these three sources are more extended than the regions of continuum emission. Although the colour and size differences quoted above are only marginally significant, these results accord with our earlier conclusion (Paper II) that the cause of the relatively broad \\lya emission lines for S1 and S3 is resonant scattering, as the photons escape though a high column density of HI. In this picture the \\lya image records the last scattering surface, which will be larger than the region of star formation. We note in passing that this implies that the observed strength of \\lya absorption and emission in spectra can depend on the slit width used.  \\subsection{The sizes of the DLA gas clouds}  In this section we make a summary of the measured impact parameters $b$ of DLA absorbers. Using this we estimate the cross-section-weighted mean radius $\\bar{R}_{DLA}$ of the gas clouds at $z>2$. Combining this information with the measured line density of DLA absorbers $dn/dz$ we are able to infer the ratio of the comoving space density of DLA absorbers at $z>2$ to the local space density of spiral galaxies.  \\begin{table*} \\begin{flushleft} \\caption[]{Measured impact parameters $b$ of DLA absorbers} \\begin{tabular}{lcccccc} \\hline\\noalign{\\smallskip} Quasar&$z_{abs}$&log(N(HI))&$b$&$b$& confirmed & reference \\\\ &  &  &($q_{\\circ}=0.0$)&($q_{\\circ}=0.5$)& & \\\\ &         &cm$^{-2}$ &kpc&kpc& & \\\\  \\hline $0454+0393$ & 0.86 & 20.8  & $4.1h^{-1}$& $3.3h^{-1}$ & N & 1  \\\\ $0302-223$  & 1.01 &$\\leq20.0\\:\\:\\:\\:\\:$& $6.2h^{-1}$& $4.9h^{-1}$ & N & 1  \\\\ $1331+170$  & 1.78 & 21.2  & $4.7h^{-1}$& $3.1h^{-1}$ & N & 1  \\\\ $0151+048$A & 1.93 & 20.4  & $7.7h^{-1}$& $5.0h^{-1}$ & Y & 2  \\\\ $1215+333$  & 2.00 & 21.0  & $8.4h^{-1}$& $5.3h^{-1}$ & N & 3  \\\\ $0841+129^a$& 2.37 & 21.3  & $8.0h^{-1}$& $4.7h^{-1}$ & N & 3,4  \\\\ $0841+129^a$& 2.48 & 21.0  & $8.0h^{-1}$& $4.7h^{-1}$ & N & 3,4  \\\\ $0528-250^b$& 2.81 & 21.3  & $8.1h^{-1}$& $4.5h^{-1}$ & Y & 5  \\\\ $2231+131$  & 3.15 & 20.0&$15.7h^{-1}\\:\\,$& $8.2h^{-1}$ & Y & 5  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\\\ $^a$ the detected galaxy could be the counterpart to one or other of the two DLA absorbers listed \\\\ $^b$ impact parameter taken from HST image (this paper) \\\\ References: 1. Le Brun et al. (1997), 2. Fynbo et al. (1997), 3. Arag\\'{o}n-Salamanca et al. (1996), 4. Wolfe et al. (1995), 5. M\\o ller and Warren (1993), 6. Djorgovski et al. (1996). \\end{flushleft} \\end{table*}  There have been many searches for optical counterparts of the DLA absorbers.  In Table 6 we have listed the small number of DLA absorbers, with redshifts $z>0.8$, for which likely candidate or confirmed optical counterparts have been detected.  Listed are the quasar name, the absorber redshift and column density, the measured impact parameter (i.e. the projected physical separation between the optical counterpart and the line of sight to the quasar), and whether or not the counterpart has been confirmed, i.e. the redshift of the counterpart has been measured to be the same as the redshift of the absorber.  In Figure 6 we plot impact parameter against column density for the absorbers listed in Table 6, for $q_{\\circ}=0.5$. The detection of the counterpart to a DLA absorber by broad-band imaging, which requires the digital subtraction of the quasar image, is more difficult for small impact parameters where the photon noise from the quasar image is greater.  Therefore it is very likely that the measured impact parameters of the few absorbers for which counterparts have been detected are larger than the average impact parameter for DLA absorbers. The vertical line in Figure 6 at ${\\rm log(N(HI))=20.3}$ marks the conventional lower limit of the column density of DLA absorbers.  The distribution of points in the figure is consistent with the expected anticorrelation between impact parameter and column density. We are interested in the mean impact parameter $\\bar{b}_{DLA}$ that would be measured if counterparts were detected for all DLA absorbers in a large unbiased survey.  From inspection of Figure 6 we suggest that it is safe to conclude that for $z>2$, $q_{\\circ}=0.5$, $\\bar{b}_{DLA}$ is less than $7h^{-1}$ kpc. For $q_{\\circ}=0.0$ we suggest $\\bar{b}_{DLA}<13h^{-1}$ kpc. The actual mean impact parameters are probably substantially smaller than these limits.  We relate the mean impact parameter to the cross-section-weighted mean radius $\\bar{R}_{DLA}$ of the absorbers by $\\bar{b}_{DLA}=\\alpha\\bar{R}_{DLA}$. The edge of the DLA is the point at which the column density falls below $N(HI)=2\\times10^{20}$ cm$^{-2}$, and we effectively assume that the column density falls off sharply at larger radii.  For face on disks $\\alpha=\\frac{2}{3}$, and for a disk that is nearly edge on $\\alpha=\\frac{4}{3\\pi}$, so we take $\\alpha=0.55$ as an average value for randomly inclined disks. On this basis we infer the following limits on the cross-section-weighted mean radius of DLA absorbers at high redshift: $$\\bar{R}_{DLA}<23.6h^{-1} {\\rm kpc}\\:\\:\\:\\: (z>2, q_{\\circ}=0.0) $$ $$\\bar{R}_{DLA}<12.7h^{-1} {\\rm kpc}\\:\\:\\:\\: (z>2, q_{\\circ}=0.5). $$  Using these limits and the measured line density of absorbers $dn/dz$ we can compute limits to the ratio of the comoving space density of DLA absorbers to the the local space density of spiral galaxies. We follow essentially the methodology employed by Wolfe et al. (1986), and Lanzetta et al. (1991, hereafter LTW). For spiral galaxies locally they adopted a galaxy luminosity function $\\Phi(\\frac{L}{L_*})$ of Schechter form, with power-law index $s$, a power-law (Holmberg) relation between radius and luminosity, of index $t$, and a ratio between gas radius and optical (Holmberg) radius $\\xi$, independent of luminosity. They found that the incidence of DLA absorbers per unit redshift $dn/dz$ at $z\\sim 2.5$ was considerably higher than expected, by a factor $F\\sim 5$, on the basis of no evolution in galaxy cross section or luminosity function normalisation. We now allow for evolution by supposing that the space density of DLA absorbers is higher than the local space density of spirals by a factor $E_{\\Phi}(z)$, and that the gas radii of galaxies are larger at high redshift by a factor $E_r(z)$. Since $dn/dz$ is proportional to the product of the space density and the galaxy cross section $\\sigma$, we have that $F=E_{\\Phi}(z)E^2_r(z)$. By comparing the expected value of the cross-section-weighted radius $\\bar{R}$ to the measured limits to $\\bar{R}_{DLA}$, we obtain limits to $E_r(z)$. Then from the measured values of $F$ we determine limits to $E_{\\Phi}(z)$, which is our goal.  Under the above assumptions the cross-section-weighted average radius of DLA absorbers is given by: $$\\bar{R}(z)=\\frac{\\int_0^{\\infty} R(\\frac{L}{L_*})\\sigma(\\frac{L}{L_*}) \\Phi(\\frac{L}{L_*})d(\\frac{L}{L_*})} {\\int_0^{\\infty}\\sigma(\\frac{L}{L_*})\\Phi(\\frac{L}{L_*})d(\\frac{L}{L_*})}$$ $$=\\frac{E_r(z)\\xi R_*\\Gamma(1+3t-s)}{\\Gamma(1+2t-s)}$$ where $R_*$ is the optical radius of a local $L_*$ spiral galaxy.  Following LTW we adopt the following values of the parameters: $t=0.4$, $s=1.25$, $\\xi=1.5$, $R_*=11.5h^{-1}$ kpc. This leads to $\\bar{R}(z)=11.0h^{-1}E_r(z)$.  Comparing with the above measured limits to $\\bar{R}_{DLA}$ we obtain the following limits to the growth factor of galaxy disks:  $$E_r(z)<2.15\\:\\: (z>2, q_{\\circ}=0.0)$$ $$E_r(z)<1.16\\:\\: (z>2, q_{\\circ}=0.5)$$  For their sample D2, of which at least 30 out of the 38 candidate DLA absorbers have been confirmed, LTW found $F=3.8$, for $q_{\\circ}=0.0$, and $F=7.1$, for $q_{\\circ}=0.5$. These results then imply that the ratio of the comoving space density of DLA absorbers at high redshift to the local space density of spiral galaxies is given by: $$E_{\\Phi}(z)=\\Phi_*(DLA)/\\Phi_*(spiral)>0.8\\:\\: (z>2, q_{\\circ}=0.0) $$ $$E_{\\Phi}(z)=\\Phi_*(DLA)/\\Phi_*(spiral)>5\\:\\: (z>2, q_{\\circ}=0.5) $$ Because the actual average impact parameters of DLA absorbers are very likely substantially smaller than the quoted limits, the actual comoving space density ratios are probably considerably greater than the limits quoted above.  \\begin{figure} \\vspace{11.5cm} \\special{psfile=Fig_bvsN.ps angle=0 voffset=-60 hoffset=-35 vscale=50 hscale=50} \\caption[ ]{Relation between impact parameter and column density for all confirmed or candidate counterparts of DLA absorbers of redshift $z>0.8$ (listed in Table 6). Different symbols correspond to different redshift ranges: squares $z>2.5$; triangles $2.5\\geq z>1.5$; circles $1.5\\geq z>0.8$. Absorbers for which the counterpart has been confirmed by spectroscopy are shown as filled symbols, and are labeled. The two possible DLA absorbers for the quasar $0841+129$ are joined by a dotted line. The hatched area shows the densely populated region from the simulation of Katz et al. (1996), and the dashed lines are the boundaries enclosing all the points in the simulation.} \\end{figure}  In Figure 5 we show also the relation between impact parameter and column density measured by Katz et al (1996) from a hydrodynamic simulation of a CDM universe with $q_{\\circ}=0.5$. Given the limited spatial resolution of the simulation, and the fact that star formation and consequent feedback were not treated, it would be premature to draw any conclusions about the apparent good agreement between the results of the simulation and the observations. Nevertheless this plot and the above calculation underscore the importance of measuring impact parameters for a large sample of DLA absorbers, a goal we are pursuing by imaging with the STIS instrument on HST.  \\subsection{The structure of the DLA absorbers}  The observations summarised above provide some indication of the typical structure of a DLA absorber. The following sketch is suggested. At the centre of the gas cloud, corresponding in projection to the highest column densities, there may be a region of star formation, of diameter $\\sim 1h^{-1}$ kpc. This central source would be observed as a Lyman-break galaxy. The gas column density decreases outwards, as suggested by the anticorrelation between impact parameter and $N_{HI}$. The size of the region over which the column density is $>2\\times10^{20}$cm$^{-2}$ is several kpc. The region of star formation would be surrounded by a zone of ionised gas. The \\lya photons are resonantly scattered in escaping to the surface of the surrounding cloud of neutral gas. The size of the observed region of \\lya emission is larger than the region of star formation, but smaller than the diameter of the gas cloud. The latter indicates a preferred direction of escape, implying that the gas resides in a flattened structure. In reality the gas cloud is likely to be irregular rather than smooth, and might contain several knots of star formation, surrounded by HII regions of different sizes, and the \\lya photons will escape preferentially where the column density of neutral gas is lowest, possibly along a complex network of tunnels.  Merging clouds will be observed as galaxies with irregular structure.   \\subsection{Filamentary structure in the distribution of galaxies at $z=3$}  We have previously suggested (Paper II) that the approximately linear arrangement of S1, S2, S3, as well as other known groups at high-redshift, may correspond to the filamentary arrangements of galaxies and galaxy sub-units that are found in computer simulations of the high-redshift universe. Figure 6 summarises the observational situation, and shows the spatial arrangement of four high-redshift groups of galaxies discovered in \\lya searches (from top to bottom: Francis et al., 1995, considering the \\lya absorber as a galaxy; Pascarelle et al., 1996; Paper I; Le F\\`{e}vre et al., 1996). All groups have been rotated to a horizontal baseline. Clearly each of the structures is elongated.  \\begin{figure} \\vspace{9.5cm} \\special{psfile=Fig_align.ps angle=0 voffset=-145 hoffset=-78 vscale=58 hscale=58} \\caption[ ]{Relative positions of galaxies in four high-redshift groups discovered in \\lya searches. Each group has been rotated so that the principal axis is horizontal in the figure, and the groups have been offset from one another vertically. Note the different scale of the lower plot.} \\end{figure}  To quantify the degree of alignment of each group we have measured the smallest internal angle $\\delta_L$ of the triangle defined by the three objects in each group (for the PKS 0528-250 field at z=2.81 we used the position of S1 for the DLA absorber, rather than the position of the quasar sightline through the absorber). The measured angles are, from top to bottom in the figure, $\\delta_L = 7.0^{\\circ}, 1.0^{\\circ}, 5.8^{\\circ}, 2.1^{\\circ}$. The average value of the angle $\\delta_L = 4.0^{\\circ}$ is certainly very suggestive of filamentariness in the distribution of galaxies at high redshift. It is also notable that the sizes of the groups range over an order of magnitude. In the computer simulations filaments are seen at all scales (e.g. Evrard et al. 1994). We do not suggest that every high-redshift group of three or four galaxies discovered will exhibit a similar degree of alignment, but rather that filamentary networks akin to those seen in the simulations may become visible as larger samples of high-redshift galaxies with denser sampling become available. The phenomenon might provide a useful discriminant for models of structure formation."
    },
    "9804/astro-ph9804033_arXiv.txt": {
        "abstract": "Recent observations supported by theoretical models have lead to the view that giant and supergiant stars are over abundant, and/or a high metallicity component may be present, in the stellar populations at the centres of active galaxies. Here we attempt to quantify these effects by observing the strengths of the stellar absorption lines of Mg~b, NaI, CaII triplet as well as molecular bands such as CN and TiO. Using long-slit spectroscopic data we are able to separate the stellar populations in and around the nucleus, for a sample including, normal, LINER, starburst and Seyfert galaxies.  In this paper we present the data, namely spectra of the nucleus and of a number of circum-nuclear regions. Comparisons reveal gradients in both the reddening and the stellar population within the central regions of most galaxies. Detailed stellar population synthesis is presented in a companion paper. ",
        "introduction": "A crucial unsolved question is whether the stellar populations in the nuclear regions of Active Galaxies differ from those of non-active galaxies of the same Hubble type. Correlations between near IR CO indices, far-infrared and X-ray luminosities of active galaxies may indicate that the more powerful monsters reside in more actively star-forming host galaxies (Yamada, 1994). The accelerated star formation caused by dynamic instabilities which trigger and/or fuel the nuclear activity could result in an overabundance of giant, supergiant and super metal rich (SMR) stars (Scoville, 1992 and references therein).  Terlevich et al. (1990) observed that in some Active Galactic Nuclei (AGN) the near IR CaII triplet absorption features are as strong, or even stronger than those of normal non-active galactic nuclei. They have suggested that the ``featureless\" blue continuum previously thought to be non-stellar in origin, actually arises from the unresolved continuum from a young cluster of stars containing red supergiants. The presence of a ``truly'' featureless continuum is then required only in the case of some Seyfert~1 galaxies which show dilution of the CaII triplet lines. Based on independent evidence we know that in some galaxies a starburst region surrounds the unresolved nucleus which emits the broad lines e.g. NGC~7469 (Wilson et al., 1991). Nonetheless this does not prove a causal connection between the starburst and AGN.  An alternative interpretation, plausible within the framework of our knowledge of AGN, is that the stellar lines are superimposed on a non-stellar nuclear continuum, and are formed in a super metal rich population. Indeed, one could expect abundance anomalies due to the intense star formation in the metal rich environment of the nucleus.  From our detailed study of the Seyfert~1 nucleus in the galaxy NGC~3516 (Serote Roos et al., 1996), we find that the stellar population exhibits a noticeable dilution by a featureless continuum in the wavelength range 5000-9800\\AA\\ {\\it as well as} a high proportion of super metal rich stars.  Thus, before dismissing the possibility of the presence of any non-thermal component in the near infrared continuum of Seyfert galaxies, we must first define the stellar population in the nucleus and the surrounding regions for a sample of AGN of all levels of activity.  Until now most studies tackling this subject have made use of only a small spectral domain around a few spectral features. In order to make further progress it is necessary to extend these studies to cover more stellar absorption features, providing signatures of different stellar populations e.g. MgI $\\lambda$5175, NaI~D $\\lambda$5896 and 8196\\AA\\ and the numerous TiO and CN bands. The NaI lines, for example, help to distinguish between the effects of overmetallicity rather than a supergiant dominated population. Indeed, the ambiguity between metallicity and supergiant excess can be resolved by comparison of the CaII triplet and NaI absorption strengths (e.g. Zhou, 1991) although the interstellar contribution to NaI introduces some uncertainty, and also by comparison with the strength of MgI (Couture \\& Hardy, 1990).  In this paper we present long-slit spectroscopy, in the range 5000-9800\\AA\\, of a sample of galaxies with different levels of activity. These observations are used to estimate radial gradients in the stellar population and to extract the stellar spectrum of the very nucleus.  The sample is selected to include various classes of active galaxies, i.e. Seyferts 1 and 2, LINERs and starbursts, in order that the strength of the stellar features may be correlated with the general properties of each class.  We also present data for two normal galaxies which will be used as comparison templates. This is not a complete sample in any statistical sense, but it does provide insights into the diversity of bulge stellar populations in and around AGN of different classes of activity.  The detailed analysis of the stellar populations for this sample using a new spectral synthesis programme (Pelat, 1997), is to be found in Paper II (Serote Roos et al., 1997). ",
        "conclusions": "The nuclei of our sample galaxies are generally redder than the outer regions (this excludes Seyfert 1s as their nuclear spectra are dominated by strong broad emission lines).  This colour gradient could either be due to dust or to stellar population gradients. In addition to dust/population gradients, the presence of a featureless continuum is inferred in a number of AGN.  This component might be of non-thermal origin, plausible in Seyfert galaxies, or of stellar origin in the case of nuclear starburst activity. The extra component dilutes the strengths of the stellar absorption lines in the spectra of the nuclear regions of AGN. A detailed study of the spectral shape of the diluting component will allow us to differentiate between the two hypotheses.  A full population synthesis analysis for most of the galaxies in our sample, using a new method developed to determine unique solutions, is presented in Paper II."
    },
    "9804/astro-ph9804080_arXiv.txt": {
        "abstract": " ",
        "introduction": "Ascertaining the core collapse supernova mechanism is a long-standing problem in astrophysics. The current paradigm begins with the collapse of a massive star's iron core and the generation of an outwardly propagating shock wave that results from core rebound. Because of nuclear dissociation and neutrino losses, the shock stagnates. This sets the stage for a shock reheating mechanism whereby neutrino energy deposition via electron neutrino and antineutrino absorption on nucleons behind the shock reenergizes it (Bethe \\& Wilson 1985; Wilson 1985).  The shock reheating phase is essential to the supernova's success, but it is precisely this phase that is difficult to simulate realistically. During shock reheating, core electron neutrinos and antineutrinos are radiated from their respective neutrinospheres, and a small fraction of this radiated energy is absorbed in the exterior shocked mantle. The shock reheating depends sensitively on the electron neutrino and antineutrino luminosities, spectra (best characterized by the {\\small RMS} energies), and angular distributions in the region behind the shock (e.g., see Burrows \\& Goshy 1993, Janka \\& M\\\"{u}ller 1996, Mezzacappa \\etal 1998).  These, in turn, depend on the neutrino transport in the semitransparent region encompassing the neutrinospheres, necessitating a neutrino transport treatment that is able to transit accurately and seamlessly between neutrino-thick and neutrino-thin regions.  Various neutrino transport approximations have been implemented in simulating core collapse supernovae. The most sophisticated approximation, which naturally has been used in realistic one-dimensional simulations, is multigroup flux-limited diffusion ({\\small MGFLD}; e.g., Bowers \\& Wilson 1982, Bruenn 1985, Myra \\etal 1987). {\\small MGFLD} closes the neutrino radiation hydrodynamics hierarchy of equations at the level of the first moment (the neutrino flux) by imposing a relationship between the flux and the gradient of the neutrino energy density (the zeroth moment).  For example,  \\begin{equation} F_{\\nu}=-\\frac{c\\Lambda}{3}\\frac{\\partial U_{\\nu}}{\\partial r}+..., \\label{eq:mgfld} \\end{equation}  \\begin{equation} \\Lambda = \\frac{1}{1/\\lambda + |\\partial U_{\\nu}/\\partial r|/3U_{\\nu}}, \\label{eq:lambda} \\end{equation} \\smallskip  \\noindent where $\\lambda$ is the neutrino mean free path, and $U_{\\nu}$ and $F_{\\nu}$ are the neutrino energy density and flux (Bruenn 1985). [Other forms for the flux-limiter $\\Lambda$ can be found in Bowers \\& Wilson (1982), Levermore \\& Pomraning (1981), and Myra \\etal (1987).] Whereas the limits $\\lambda \\rightarrow 0$ and $\\lambda \\rightarrow \\infty$  produce the correct diffusion and free streaming fluxes, it is in the critical intermediate regime where the {\\small MGFLD} approximation is of unknown accuracy. Unfortunately, the quantities central to the postshock neutrino heating, i.e., the neutrino {\\small RMS} energies, luminosities, and mean inverse flux factors, are determined in this regime, and given the sensitivity of the neutrino heating to these quantities, it becomes necessary to consider more accurate transport schemes.  Moreover, in detailed one-dimensional simulations that have implemented elaborate {\\small MGFLD} neutrino transport (e.g., see Bruenn 1993, Wilson \\& Mayle 1993, and Swesty \\& Lattimer 1994), explosions were not obtained unless the neutrino heating was boosted by additional phenomena, such as convection. This leaves us with at least two possibilities to consider:  (1) Failures to produce explosions in the absence of additional phenomena, such as convection, have resulted from neutrino transport approximation.  (2) Additional phenomena may be essential in obtaining explosions. ",
        "conclusions": "Comparing three-flavor {\\small MGBT} and three-flavor {\\small MGFLD} in postbounce supernova environments, we find that {\\small MGBT} leads to a significant increase/decrease in the {\\it net} heating/cooling rate, particularly  just above/below the gain radius. The {\\small MGBT} net heating  rate can be as much as 2 times the {\\small MGFLD} net heating rate above the gain radius, with net cooling rates that are typically 0.8 times the {\\small MGFLD} rate below. These differences stem primarily from differences in the neutrino luminosities and mean inverse flux factors;  the heating rate is linearly proportional to both these quantities, and differences in both add to produce a significant difference in the net heating rate.  We also observe that the differences in the net heating rate are greatest at earlier postbounce times for a given progenitor mass, and at a given postbounce time, greater for greater progenitor mass. This is illustrated in Table 1. The enhancement in heating with increased progenitor mass suggests that the net heating enhancement from {\\small MGBT} is potentially robust and self-regulated.  In closing, our results are promising, and their ramifications for core collapse supernovae,  and in particular, for the postbounce neutrino-heating, shock-revival mechanism, await one- and two-dimensional dynamical simulations with {\\small MGBT} coupled to the core hydrodynamics. One-dimensional simulations are currently underway, and we plan to report on them soon.  \\begin{table}[t] \\vspace{-24pt} \\begin{center} \\footnotesize\\rm \\caption{Maximum Net Heating/Cooling Rates \\label{t1}} \\begin{tabular}{lccc} \\topline {Progenitor Mass [${\\rm M}_{\\odot}$]} & {$t_{pb}$ [ms]} & {Maximum Net Heating Ratio} & {Maximum Net Cooling Ratio} \\\\ \\midline 15& 106& 2.0& 0.8 \\\\ & 233& 1.3& 0.8 \\\\ 25 & 156& 2.0& 0.8 \\\\ \\bottomline \\end{tabular} \\end{center} \\vspace{-12pt} \\end{table}"
    },
    "9804/astro-ph9804049_arXiv.txt": {
        "abstract": "We investigate in detail the hypothesis that low surface brightness galaxies (LSB) differ from ordinary galaxies simply because they form in halos with large spin parameters.  We compute star formation rates using the Schmidt law, assuming the same gas infall dependence on surface density as used in models of the Milky Way.  We build stellar population models, predicting colours, spectra, and chemical abundances.  We compare our predictions with observed values of metallicity and colours for LSB galaxies and find excellent agreement with {\\it all} observables. In particular, integrated colours, colour gradients, surface brightness and metallicity match very well to the observed values of LSBs for models with ages larger than 7 Gyr and high values ($\\lambda > 0.05$) for the spin parameter of the halos.  We also compute the global star formation rate (SFR) in the Universe due to LSBs and show that it has a flatter evolution with redshift than the corresponding SFR for normal discs.  We furthermore compare the evolution in redshift of $[Zn/H]$ for our models to those observed in Damped Lyman $\\alpha$ systems by \\scite{Pettini+97} and show that Damped Lyman $\\alpha$ systems abundances are consistent with the predicted abundances at different radii for LSBs. Finally, we show how the required late redshift of collapse of the halo may constrain the power spectrum of fluctuations. ",
        "introduction": "Late type low surface brightness galaxies are found to be a significant fraction of the total galaxy population in the Virgo cluster (\\pcite{Impey+88}), in the Fornax cluster (\\pcite{Irwin+90}), and in the field (\\pcite{Mcgaugh+95}; \\pcite{Sprayberry+96}; \\pcite{Sprayberry+97}).  A theory of galaxy formation should therefore account for the existence of LSBs.  On the other hand, LSBs are particularly useful because they are simpler objects than high surface brightness (HSB) galaxies: they have relatively little present day star formation and little dust (reddening). Therefore it is easier to model their stellar population and star formation history.  They are also more dark matter dominated than HSBs, and therefore particularly suitable for probing the structure of dark matter halos.  In order to compare theories of galaxy formation with observational surveys of galaxies it is important not only to quantify observationally the abundance of LSBs, but also to understand their nature. For instance, if LSBs are made of very young stellar populations, as often claimed in the literature, they may be irrelevant in the picture of the Universe at high redshift, while they could be a substantial component of that picture if their stellar populations are old.  The existence of LSB galaxies can be readily understood if it is assumed, following \\scite{Fall+80}, that the specific angular momentum of the baryons is approximately conserved during their dissipation into a rotationally supported disk, and that the disk length--scale is therefore related to the angular momentum of the dark matter halo. The low surface brightness is the consequence of the low surface density of the disk, which is due to the larger spin parameter of the dark matter halos of LSBs relative to the spin parameter of the halos of HSB disks (e.g. \\pcite{Dalcanton_disc+97}; \\pcite{Mo+97}; \\pcite{Jimenez+Heavens+97}).  However, while the surface density of a disk can be simply related to the mass and spin parameter of its dark matter halo, its surface brightness, and therefore its mass to light ratio ($M/L$), depends on the stellar population and could in principle vary with galaxy type and age. For instance, $M/L$ must have a significant time dependence due to the history of star formation in the disk and to the continuous death of massive stars. Moreover, $M/L$ in a given photometric band changes with time also as a consequence of the colour evolution of the stellar population, which is sensitive also to the history of the chemical enrichment of the disk.  Therefore, a comparison between a simple model of disk formation and the observations cannot be made without an appropriate model for the stellar population in the disk that includes {\\it self-consistently} the chemical evolution of the population.  \\begin{figure} \\centering \\leavevmode \\epsfxsize=1. \\columnwidth \\epsfbox{surf_1p0.eps} \\caption[]{Initial surface density for an isothermal sphere and 3 different masses. The solid lines correspond to $\\lambda=0.1$ to 0.01 (from top to bottom) in steps of 0.01. The circular velocity of the halo (dashed line) is also plotted.} \\label{f1} \\end{figure}  In this paper we model self-consistently the chemical and photometric evolution of the stellar population of LSB disks, formed in dark matter halos described by isothermal spheres.  The aim of this paper is twofold:  \\begin{enumerate} \\item show that only one parameter, namely the {\\it spin parameter $\\lambda$} of the dark halo, can explain the surface brightness of LSBs. \\item use LSB disk galaxies to study the formation and evolution of all disk galaxies. \\end{enumerate}  The second point is motivated by the first, that is by the fact that LSB disks are in fact so similar to HSB disks. On the other hand, LSBs are more dark matter dominated than HSBs, and therefore better described by a very simple model where only the gravitational field of an isothermal halo is considered.  Our assumptions are as follows:  \\begin{enumerate} \\item{The specific angular momentum of baryonic matter is the same as for dark matter} \\item{Gas settles until centrifugally supported, in a given dark matter halo} \\item{The star formation rate is given by the Schmidt law} \\item{The gas infall rate is assumed to be the same function of total surface density as used in models of the Milky Way} \\end{enumerate}  \\begin{figure*} \\centering \\leavevmode \\epsfxsize=1.6 \\columnwidth \\epsfbox{surf_o.eps} \\caption[]{Contour plot of the value of the central surface density (in M$_{\\odot}$pc$^{-2}$) for different values of the halo mass and spin parameter. Left panel: halo formation redshift $z=0.7$. Right panel: halo formation redshift $z=2.1$.} \\label{f2} \\end{figure*}  We first use the density profile of the halo to compute the surface density of the settling disk. We then use the Schmidt formation law and an infall rate that reproduces the observed properties of the Galaxy in conjunction with the chemical evolution models by \\scite{Matteucci+89} to compute the star formation rate at several radii of the disk and the evolution of several chemical species (H, D, He, C, N, O, Ne, Mg, Si, S, Ca, Fe and Zn). We then proceed to compute spectra, integrated colours, colour profiles and surface brightness for two different values of the spin parameter of the halo.  We first present in the next section the simple non--self--gravitational disk model and the non--singular halo model. We describe in section 3 the set of synthetic stellar population models used to predict the spectra and colours of the stellar population.  In section 4 we model the star formation in the disk with a Schmidt law and a gas infall law, assuming the same model of the chemical evolution of the Galaxy as \\scite{Matteucci+89}. Finally, in section 5 we compute synthetic spectra and photometric properties of the stellar population in the disk, using the stellar evolution tracks of JMSTAR15, stellar atmosphere models by \\scite{Kurucz_92} and J{\\o}rgensen (private communication) and the chemical evolution models built for LSBs (see section 4).  We discuss the results in section 6.  The most important results of this work are:  \\begin{enumerate} \\item Observed colour profiles, chemical abundances and surface brightness profiles for LSBs are well fitted if they are assumed to have a spin parameter for its halo higher than HSBs. \\item LSBs are not young objects, as often claimed in the literature, since their colours are well fitted by old ($> 7$ Gyr) stellar populations. \\item There is discrepancy between the photometric age of the galaxies and the age of formation of their halo, which indicates that the star formation can start about 2~Gyr before the halo is formed. This discrepancy is reduced to 1~Gyr if the Hubble constant is assumed to be $H_0=65$~kms$^{-1}$Mpc$^{-1}$ and completely removed if the Universe is open ($\\Omega < 0.3$) or has a significant vacuum energy contribution ($\\Lambda > 0.6$). \\item The earliest stellar populations of LSBs could be present in the high redshift universe in a similar proportion to HSBs as they are now. Their colours at high redshift (as the colours of all disks) are about 1~mag bluer than at low redshift. \\end{enumerate}  In a previous paper (\\pcite{Padoan+Jimenez+Antonuccio97a}) we have shown that LSBs are not necessarily young and un-evolved systems (e.g. \\pcite{McGaugh+Bothun94,deBlok_phot+95}).  We used the simplest possible model, a burst of star formation, to determine a lower limit to the age of LSBs studied by \\scite{deBlok_phot+95}. In the present work we considerably improve our previous model because we use a more realistic continuous star formation process, we include self--consistently the chemical evolution, and we connect the disk model to the cosmological scenario using the spherical collapse model (\\pcite{Gunn_Gott_72}).  This more detailed model confirms our previous result that the blue LSB disks in the sample by \\scite{deBlok_phot+95} are not un-evolved objects collapsed at late times from low initial over-densities (\\pcite{McGaugh+Bothun94,Mo+94}), but rather normally evolved disk galaxies.  Throughout the paper the present day value of the Hubble constant is assumed to be $H_0=75$~kms$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\subsection{The Nature of LSBs}  \\begin{figure*} \\centering \\leavevmode \\epsfxsize=1.4 \\columnwidth \\epsfbox{lsbabund.eps} \\caption[]{The redshift evolution for some elements is plotted. The top-left panel shows the evolution of $[O/H]$ and is compared with \\scite{McGaugh_oxigen94} measures. Our model reproduces perfectly the peak in $[O/H]$ among LSBs found by \\scite{McGaugh_oxigen94} at $[O/H]=-0.9$. The top-right panel shows the analogous to the previous one but for $[Zn/H]$ compared with \\scite{Pettini+97} measures. It transpires from the figure that the predicted $[Zn/H]$ for LSBs is consistent with the spread found for DLAs.} \\label{f17} \\end{figure*}  In this work we have shown that the photometric properties of the bluest galaxies in the sample by \\scite{deBlok_phot+95} can be reproduced assuming that they form and evolve as normal disk galaxies, with relatively high spin parameter. That the spin parameter alone could explain the low surface brightness was already shown by \\scite{Dalcanton_disc+97}, but they assumed a given $M/L$ ratio. Here we have strengthened the point by also explaining the colours, the colour gradients, and the chemical abundances. We have also provided synthetic spectra that could be compared with the observations, in the effort to prove that LSB disk galaxies are indeed normal galaxies with high spin parameter.  The stellar populations of LSBs are rather evolved, especially in the central part of their disks. Even the bluest galaxies in the LSB sample by \\scite{deBlok_phot+95} are older than 7~Gyr, and typically star formation in these galaxies starts about 9~Gyr ago. This confirms the results of \\scite{Padoan+Jimenez+Antonuccio97a}, where the galaxies in the same sample were found to be at least 7~Gyr old.  LSB and HSB disks are therefore hosted in similar dark matter halos that differ only for their spin parameter, and they evolve in a similar way, apart from the fact that LSB disks are less concentrated than HSB disks. This means that the present day (or low redshift) abundance of LSBs relative to HSBs should be about the same at any redshift. Moreover, the conclusions about the epoch of star formation and halo formation (see the next section), and about the colour evolution must be valid for disk galaxies in general.  In particular, Fig.~\\ref{f11} shows that the colour evolution of LSBs is very strong. The U-B, B-V and V-I colours of the central part of the disk at high redshift are about 1~mag bluer than they are at low redshift, and the B-R colour even 1.5~mag bluer.  This must be a property of disks in general, and not only of LSB disks.  \\subsection{Star Formation and Galaxy Formation}  We have shown in section 2.1 that the size of LSB disk galaxies is such that their halos are formed in the redshift range $0.1<z<1.2$. We have also been able to estimate the formation redshift of the halo of each galaxy in the sample, and we have found a mean value $z=0.7$.  In an Einstein--de--Sitter universe with $H_0=75$~kms$^{-1}$Mpc$^{-1}$, a redshift $z=0.7$ corresponds to a look--back time of $6.8$~Gyr.  On the other hand, the photometric properties of LSB galaxies, in particular their colours, are such that star formation must have started approximately $9$~Gyr. Therefore, the star formation starts about 2~Gyr before the galactic halos are formed. These numbers are likely to be similar for HSB disk galaxies too, since they have colours and sizes very similar to LSBs, also suggesting a rather low formation redshift for their halos.  The discrepancy between the epoch when star formation starts in a galaxy and the epoch when the halo forms means that the process of star formation starts before the galactic halo has assembled, as predicted for example by bottom--up galaxy formation scenarios, characterized by power spectra with more power on sub--galactic scales than on galactic ones.  In Fig.~\\ref{f18} we show the evolution with redshift of the SFR. The upper panel shows the fraction of gas with time that is converted into stars at every redshift for the nucleus and the disk. The continuous thick line shows the average of the nucleus and the disk as a representative of the whole LSB. It is worth noticing that most of the gas in the nucleus is converted into stars between $2 < z < 4$, while most of the star formation in the disk takes place at $ z < 2$. The lower panel shows the global SFR in the Universe for LSBs assuming that they have the same comoving number density as HSBs and that the typical mass for a $L_*$ LSB is a factor 5 smaller than the typical mass of a normal $L_*$ galaxy. The data points are taken from \\scite{Madau_97} and trace the global SFR in the Universe for HSBs. LSBs lie below the observed points because they have typical masses smaller than HSBs. The important point to notice here is that the global SFR for LSBs has a flatter evolution than the observed points.  \\begin{figure} \\centering \\leavevmode \\epsfxsize=1.0 \\columnwidth \\epsfbox{sfr.eps} \\caption[]{Top panel shows the gas fraction with time in our LSB model that is converted into stars $vs$ redshift for the disk and the nucleus. The thick line corresponds to the average of the nucleus and the disk and should be a fair representative of the whole LSB. The bottom panel shows the predicted global SFR in the Universe due to LSBs assuming that the mass for a $L_*$ LSB is 5 smaller than the mass for a $L_*$ HSB and that the comoving number density for LSBs is the same as for HSBs (see text). The observed points are taken from \\scite{Madau_97}. Notice that the evolution with redshift is flatter than the one measured by \\scite{Madau_97}.} \\label{f18} \\end{figure}  \\subsection{Discussion}  In order to explain the rotational properties of LSBs, we require halos with velocity dispersions $\\approx 100$ km s$^{-1}$, or masses $\\approx 10^{11}$ M$_\\odot$.  This is also consistent with the required gas masses and a baryon fraction of a few percent.  To obtain the right sizes of LSBs, we additionally require a relatively late halo collapse, $z \\simeq 0.7$.  Clearly this may present a problem for galaxy formation models which have a lot of power on the appropriate scale.  An obvious example is COBE-normalised CDM, with shape parameter $\\Gamma=0.5$ \\cite{FW91}, for which the characteristic collapse redshift for $100$ km s$^{-1}$ halos is around $z=6$.  Reducing the amplitude to that of standard CDM (bias parameter 2.5) alleviates the problem to a certain extent, but the collapse redshift is still $z=3$.  It is more promising to consider mixed dark matter models, for which the power spectrum is in any case a better match to the observed galaxy spectrum.  For a neutrino fraction of 30\\%, the comoving number density of halos with mass exceeding $1.5 \\times 10^{11}$ M$_\\odot$ reaches 0.01 $h^3$ Mpc$^{-3}$, comparable to the density of large galaxies, at a redshift of about 1.7 \\scite{KBHP95}. This is still a little high, but it should be remembered that the collapse of halos of given mass does take place over a reasonably wide range of redshifts, and some fraction of halos would form after $z=0.7$.  In the model considered by \\scite{KBHP95}, about a quarter of halos exceeding this mass form after $z=0.7$.  Analytically, there is expected to be a weak correlation between high spin and late collapse \\scite{HP88}, and numerical simulations by \\scite{Ueda94} show an anticorrelation between spin and density.  We postulate therefore that the LSBs form in halos with high spin which form relatively late, in the context of a hierarchical model.  The indications are that a model with relatively little small-scale power, such as mixed dark matter, would do best here, but there are enough uncertainties in the collapse redshift, both in the simplifications of the halo model, and in physical processes which might delay collapse (e.g. \\pcite{BabulRees92}) that this conclusion is not strong.  In summary, the most important results of this work are:  \\begin{enumerate} \\item Observed colour profiles, chemical abundances and surface brightness profiles for LSBs are well fitted if they are assumed to have a spin parameter for its halo higher than HSBs. \\item LSBs are not young objects, as often claimed in the literature, since their colours are well fitted by old ($> 7$ Gyr) stellar populations. \\item There is discrepancy between the photometric age of the galaxies and the age of formation of their halo, which indicates that star formation can start about 2~Gyr before the halo is formed.  This is perfectly acceptable in hierarchical models of galaxy formation, and, as discussed in the text, the discrepancy can be reduced or removed if the cosmological model we have assumed is incorrect. If an open ($\\Lambda$) cosmology had been adopted, the above discrepancy would have been completely removed. \\end{enumerate}"
    },
    "9804/astro-ph9804339_arXiv.txt": {
        "abstract": "The variable star population of the galactic globular cluster NGC 1851(C0512-400) has been studied by CCD photometry, from observations made in the B, V and I bands during 1993-1994.  Light curves are presented for 29 variables, seven of which are new discoveries.  The behavior of the RR Lyraes in the period-temperature diagram appears normal when compared to clusters which bracket the NGC 1851 metallicity.  Reddening and metallicity are re-evaluated, with no compelling evidence being found to change from the values of $E(B-V) = 0.02$ and $[Fe/H] = -1.29$ (Zinn scale) adopted in recent studies of the cluster.  Photometry is provided for stars in an annulus with radii 80 and 260 arcsec centered on NGC 1851.  To at least $V=18.5$ there is excellent agreement with the extensive earlier photometry for the brighter NGC 1851 stars, with systematics less than 0.02 mag in all colors.  Instability strip boundary positions for several clusters shows a trend for the red boundary to move to redder colors as the metallicity increases.  keywords: globular clusters: individual (NGC1851) - RR Lyrae variable ",
        "introduction": "The galactic globular cluster (GC) NGC 1851 (C0512-400) is rich, centrally-condensed and belongs to the small group of clusters which display bimodal horizontal-branch (HB) morphology, defined (Catelan et al 1998) as having fewer RR Lyrae stars than either blue or red HB stars.   Canonical theory, that is considering a GC as a population characterized by a single age, constant abundance and a red-giant branch (RGB) mass loss parameter that has a narrow Gaussian distribution (typically $\\sigma_M \\sim 0.02 M_{\\sun}$), cannot explain such unusual HB morphology.   In order to account for the bimodality, attention has focussed recently on scenarios which can alter the mass loss parameter, such as tidal stripping of red-giant envelopes in dense environments, rapid rotation, stellar encounters, and binary interactions.  Sosin et al (1997) discuss these various options in the context of the most extreme example known of a GC with bimodal HB, NGC 2808, which displays a blue HB with multiple gaps that extends to below the main sequence turnoff in the $V, B-V$ color-magnitude diagram (CMD),  $M_V \\sim 5$.  They conclude that none of the present explanations are a satisfactory match to the observations.   However, Sweigart \\& Catelan (1997) have modeled the unusual HB morphology of the metal rich GC's NGC 6338 ($[Fe/H] = -0.60$) and NGC 6441 ($[Fe/H = -0.53$) for which Rich et al (1997) have obtained CMD's using the Hubble Space Telescope.  Both clusters have HB's which slope upwards (brighter) with decreasing $B-V$, and have extended blue tails.  Models with high helium abundance, rapid rotation, and helium mixing into the envelope are all able to produce a sloping HB morphology, and sometimes a bimodal distribution.  The helium mixing alternative is particularly interesting given the observed heavy-element abundance variations in globular cluster red-giant stars (Kraft 1994).  Mixing deep enough to produce enhanced aluminium, as observed in some stars, will also dredge up helium.  Extensive deep mixing might be expected to destroy the sharp boundary corresponding to the deepest penetration of the convective zone, and thus prevent the observational pile-up of stars on the RGB near the level of the HB. NGC 1851 in fact appears to have quite a prominent such clump, thus suggesting that deep mixing has not taken place on the RGB at a luminosity less than that of the clump. Notwithstanding, the number of possible options still available to explain the peculiar NGC 1851 HB is considerable.   The CMD of NGC 1851 has most recently been studied by Walker (1992) (hereafter W92) in the $B$ and $V$ bands, in the UV by Parise et al (1994) and in the $V$ and $I$ bands by Saviane et al. (1997) (hereafter S97), where references to earlier work can be found. In both optical studies the bimodal HB is interpreted as a consequence of differing efficiencies of mass loss as the stars evolve up the RGB.  W92 suggested that a unimodal mass distribution might be able to produce a bimodal HB stellar distribution when the detailed shape of the evolutionary tracks was taken into account, however S97 do not find good agreement when comparing with the Bertelli et al. (1994) tracks.  They prefer a bimodal mass loss distribution, and indeed find some evidence to suggest that the radial distributions of the blue and red HB stars differ, pointing towards some as yet unexplained interaction between the dynamical evolution and the stellar evolution of these stars.  On the other hand, Catelan et al (1998), using updated (Sweigart 1997) Sweigart and Gross (1976, 1978) models, find they can reproduce the NGC 1851 HB morphology with a unimodal, albeit very wide, mass distribution having characteristics $<M_{HB}> = 0.665 M_{\\sun}$ and $\\sigma_M = 0.055 M_{\\sun}$.  Stetson et al. (1996) and Sarajedini et al. (1997) review the question of the relative ages of globular clusters, to reach very different conclusions.  In both cases the HB level of NGC 1851 (W92) is used to link the second-parameter pair NGC 288 and NGC 362.  Critical to these arguments is the V magnitudes of the reddest BHB stars and the RR Lyraes; the latter will be provided here for the first time.  Lying in the region of the HB which is sparsely populated, the RR Lyraes may provide important clues to help explain the reason for the bimodal HB, given the constraints that the pulsation properties place on stellar evolutionary status. There have been suggestions (Catelan 1997) that the 1851 variables are peculiar with respect to their behavior in the period-temperature diagram and that the photographic studies, as detailed below, also show that several of the RRab stars have light curve amplitudes near 2 magnitudes, much larger than normal.  Sawyer Hogg (1973) lists 10 variables in NGC 1851, from discoveries by Bailey (1924) and Laborde and Fourcade (1966).  Preliminary periods for some of these stars, and an additional four new discoveries, resulted from a short observing campaign by Liller (1975) using photographic photometry at the CTIO 1.0-m and 1.5-m telescopes.  She noted that V2 and V8 appeared to be constant, and V9 was very red.  Wehlau et al. (1978) (hereafter W78) measured an additional 57 plates, almost all taken with the 1.0-m Swope telescope at Las Campanas, and analysed them along with the Liller (1975) plates.  Periods were derived and light curves presented for a total of 19 RR Lyrae stars, the mean period of the RRab stars being 0.573 days, and the ratio of the number of RRc to RRab variables was found to be 0.36, both are values typical of an Oosterhoff type I system.  The photometric zeropoint calibration for these observations was very uncertain, due to the lack of a definitive photometric sequence in the field.  Stetson (1981) found four additional variables and two apparently constant stars lying in the instability strip. Wehlau et al. (1982) (hereafter W82) studied these stars using their original plate material supplemented with another 18 plates taken in 1970, whereupon three stars were found to be RR Lyraes, a fourth was classified as a probable field W UMa star, while a fifth is a red variable.  W78 and W82 determined accurate periods and approximate mean $<B>$ magnitudes for all the stars they identified as RR Lyraes, classifying 15 stars as RRab and seven as RRc. The light curves display the $0.1-0.2$ mag scatter typical for the photographic technique.  No modern studies of the NGC 1851 variables appear in the literature, however S97 identify an additional seven candidate RR Lyrae variables from a comparison of their photometry with that of W92, selecting stars with deviant photometry and $V \\sim 16, V-I \\sim 0.4$. Here we present new CCD photometry for the NGC 1851 RR Lyrae variables and compare with results for RR Lyraes in other GC's in this program (IC 4499, M68, M72, NGC 6362). ",
        "conclusions": ""
    },
    "9804/astro-ph9804027_arXiv.txt": {
        "abstract": "The Hubble Deep Field-South was chosen to have a QSO (RA 22:33:37.6 Dec $-$60:33:29 J2000 and B=17.5) in the field to allow for studies of absorption systems intersecting the sight line to the QSO.  To assist in the planning of HDF-S observations we present here a ground-based spectrum of the QSO. We measure a redshift of $z=2.24$ for the quasar and find associated absorption in the spectrum at $z=2.204$ as well as additional absorption features. ",
        "introduction": "Unlike the original Hubble Deep Field, the Hubble Deep Field South (HDF-S) was chosen specifically to contain a $z>2$ QSO suitable for studying the relationship between the high redshift galaxies identified in the HDF-S and the absorption lines in the spectrum of the HDF-S QSO. The QSO was found on a UK Schmidt Telescope objective prism plate scanned by Mike Irwin using the Automated Plate Measuring facility in Cambridge analysed by Paul Hewett and then confirmed by observations at the Anglo-Australian Telescope (Boyle 1997). To aid future observations of the HDF-S we present here a low-resolution spectrum of the QSO; the data are available at http://bat.phys.unsw.edu.au/$\\sim$kms/hdfs/. ",
        "conclusions": ""
    },
    "9804/astro-ph9804211_arXiv.txt": {
        "abstract": "We use the measures of Li and rotational velocities in F Hyades stars to assess the role of the wind-driven meridian circulation and of shear turbulence in the transport of angular momentum in stars of different masses. Our models include both element segregation and rotation-induced mixing, and we treat simultaneously the transport of matter and angular momentum as described by Zahn (1992) and Maeder (1995).  We show that the hot side of the Li dip in the Hyades is well explained within this framework, which was also successfully used to reproduce the C and N anomalies in B type stars (Talon et al. 1997). On the cool side of the dip, another mechanism must participate in the transport of angular momentum; its efficiency is linked to the depth of the surface convection zone. That mechanism should also be responsible for the Sun's flat rotation profile. ",
        "introduction": "During the last decade, special efforts have been devoted to improve the description of the mixing processes related to stellar rotation. The most recent works (see for example Pinsonneault et al. 1989, Zahn 1992, Maeder 1995, Talon \\& Zahn 1997) describe the evolution of the internal distribution of angular momentum in a self-consistent manner under the action of meridional circulation and of shear turbulence. The mixing of chemicals is then linked directly to the rotation profile, whereas previous studies made use merely of a parametric relation between the turbulent diffusivity and the rotational velocity (cf. e.g. Schatzman et al. 1981, Zahn 1983).  Such a self-consistent treatment was applied successfully by Talon et al. (1997) in the study of a 9 M$_{\\odot}$ star, modeling the transport of angular momentum by the meridional circulation as a truly advective process. The only assumption in this theory is that the turbulence sustained by the shear is highly anisotropic and relies on two free parameters; the first one describes the magnitude of the horizontal shears (cf. Zahn 1992) and the second one, the erosion of the restoring force due to both the thermal and the mean molecular weight stratifications (cf. Maeder 1995, Talon \\& Zahn 1997). These authors reproduce the slight abundance anomalies measured in B stars by Gies \\& Lambert (1992). They also show that the widening of the main sequence, which is generally attributed to convective overshooting in massive stars, may be due to the rotational mixing present in stars having a ``typical'' velocity for the spectral type considered.  Concerning low-mass stars, it has been shown that the hydrodynamical models relying on meridional circulation and shear fail to reproduce the solar rotation profile given by the helioseismic observations (Brown et al. 1989, Kosovichev et al. 1997): at the solar age, those models still have large $\\Omega$ gradients which are not present in the Sun (see Chaboyer et al. 1995 and Matias \\& Zahn 1997). That conclusion has been reached independently by two different groups, using different descriptions for the transport processes. On one hand, the Yale group computed the evolution of angular momentum in low mass stars with a simplified description of the action of the meridional circulation which was considered as a diffusive process rather than as an advective process. The whole evolution of momemtum and chemicals was then due to diffusion only, with a free parameter that had to be calibrated to differentiate the transport of the passive quantities with respect to that of vectorial ones. Pinsonneault et al. (1990) were then able to reproduce the surface Li abundances for low-mass cluster stars (with effective temperature lower than 6500K). However, they obtained large rotation gradients within these stars which are excluded by helioseismology (Chaboyer et al. 1995). On the other hand, Matias \\& Zahn (1997) performed a complete study for the evolution of the Sun's angular momentum, where they took into account the advective nature of the meridional circulation. They also concluded that meridional circulation and shear turbulence are not efficient enough to enforce the flat rotation profile measured by helioseismology.  These results indicate that another process participates in the transport of angular momentum in solar-type stars, while the so-called wind-driven meridional circulation (Zahn 1992) is successful in more massive stars. In order to study the transition between solar-type and more massive stars and to identify the mass range for which the present description for the transport of angular momentum and chemicals relying only on rotation fails, we propose to use the measures of lithium and rotational velocities in galactic cluster stars.  We first review the observations of lithium abundances and rotation in the Hyades main-sequence stars, and summarize the difficulties of the various models proposed so far to explain the Li dip in F stars (\\S 2). We recall the equations that describe the evolution of angular momentum due to meridian circulation and shear turbulence as well as the associated transport of chemicals (\\S 3). We study the impact of rotational mixing on the lithium abundance in galactic cluster F stars, and compare this to the observations. Our models include both element segregation and rotation-induced mixing, and we treat simultaneously the transport of matter and angular momentum. The internal rotation profile thus evolves completely self-consistently under the action of meridional circulation as described by Zahn (1992) (see also Matias et al. 1997), and of shear stresses which take into account the weakening effect of the thermal diffusivity, as was first shown by Townsend (1958) (\\S 4). We show that the blue side of the lithium dip is well reproduced within this framework, and that the process responsible for the shape of the solar rotation profile should become efficient only for stars on the cool side of the Li dip, where the external convection zone is thick enough. By achieving efficiently momentum transport, the global effect of this process would be to reduce the mixing due to the rotational instabilities in stars with effective temperature lower than $\\sim$ 6500K. The most likely candidates for this transport process are the gravity waves generated by the external convection zone (Schatzman 1993; Zahn et al. 1997; Kumar \\& Quataert 1997) and the large-scale magnetic field which could be present in the radiative interior (Charbonneau \\& MacGregor 1993; Barnes et al. 1997). ",
        "conclusions": "Assuming rapid rotation and using a self-consistent description for the transport of angular momentum and of chemicals by meridional circulation and shear instabilities (cf. Zahn 1992, Talon \\& Zahn 1997), Talon et al. (1997) successfully explained the C and N anomalies observed in some B stars.  At the same time, it was shown (Matias \\& Zahn 1997) that this description applied to the transport of angular momentum in the Sun is incomplete, leading to large $\\Omega$ gradients which are not observed. Another transport mechanism must thus be invoked in low mass stars.  At this point, 2 questions remain : firstly, the nature of that transport mechanism has to be determined unambiguously and secondly, the location of the transition between the regime which is relevant for massive stars and the one which is relevant for low mass stars has to be identified.  In this paper, we addressed that second question. We presented numerical calculations of Li destruction due to rotational mixing using the {\\it same} description as Talon et al. used for more massive stars and the same free parameters. We showed that this clearly reproduces the hot side of the Li dip. Let us recall that the destruction of lithium is then due solely to rotational mixing enhanced by the spin down of the outer layers. Stars hotter than 7000 K also undergo rotational mixing, but it is much milder due to the weak differential rotation.  The rise of Li abundances on the right side of the dip is not explained within this framework. We propose that it is linked to the appearance of another transport mechanism for angular momentum which reduces the magnitude of the meridional circulation and shears, leading to the observed diminution of Li destruction on the red side of the Li dip. This mechanism is known to occur in the Sun where it is responsible for the flat rotation profile."
    },
    "9804/astro-ph9804161_arXiv.txt": {
        "abstract": "The Berkeley spectrograph aboard the ORFEUS telescope made its second flight on the 14-day ORFEUS-SPAS II mission of the Space Shuttle {\\it Columbia} in November/December 1996.  Approximately half of the available observing time was dedicated to the Berkeley spectrograph, which was used by both Principal and Guest Investigators.  The spectrograph's full bandpass is 390--1218 \\AA; here we discuss its in-flight performance at far-ultraviolet (FUV) wavelengths, where most of the observations were performed.  The instrument's effective area peaks at 8.9 $\\pm$ 0.5 cm$^2$ near 1020 \\AA, and the mean spectral resolution is 95 km s$^{-1}$ FWHM for point sources.  Over most of the spectral range, the typical night-time background event rate in each spectral resolution element was about 0.003 s$^{-1}$.  Simultaneous background observations of an adjacent blank field were provided through a secondary, off-axis aperture.  The Berkeley spectrograph's unique combination of sensitivity and resolution provided valuable observations of approximately 105 distinct astronomical targets, ranging in distance from the earth's own moon to some of the brightest AGN. ",
        "introduction": "The German spacecraft Astro-SPAS, a deployable platform designed to meet the technical performance demands of astronomical payloads and similar scientific instruments, comprised the primary payload aboard shuttle mission STS-80 ({\\it Columbia}). On this, its third flight, the platform carried a trio of far-ultraviolet instruments: two independent spectrographs within the 1 meter diameter ORFEUS telescope (Grewing et al. 1991) and the IMAPS objective-grating spectrograph (Jenkins et al. 1996). All three had flown on the Astro-SPAS' 5-day maiden voyage in September of 1993, but improvements in instrument performance, and the critical need for additional observation time, motivated a reflight. A photograph of the payload is shown in Plate~1.  With few exceptions, targets suitable for ORFEUS were too faint for IMAPS, so no attempt was made to coalign these instruments closely.  Within the ORFEUS telescope a flip mirror was employed to direct the optical beam to one spectrograph or the other. Hence in general only one of the three instruments was operated at a time. The available observing time was shared equally between Guest Investigators selected by peer review and the Principal Investigator teams who had provided the instruments. Flight operations were directed from a control complex at the Kennedy Space Center.  The general design of the Berkeley spectrograph has been discussed previously (Hurwitz \\& Bowyer 1986, 1996). We changed the instrument between missions only by overcoating of two of the four diffraction gratings (including the far ultraviolet grating) with silicon carbide, introducing the multiple apertures discussed below, and modifying the detector electronics to improve the imaging at high count rates. We did not recoat the KBr photocathode on the microchannel plate detectors; the delay-line anode detector systems are discussed in Stock et al. (1993). In this work we report on the performance and calibration of the spectrograph during the ORFEUS-SPAS II mission and the instrumental effects of interest to Guest Investigators and other users of the extracted data products. ",
        "conclusions": "The Berkeley spectrograph aboard the ORFEUS telescope offers a unique and important combination of spectral resolution and effective area in the comparatively unexplored far-ultraviolet wavelength band. During the ORFEUS-SPAS II mission in November/December 1996, Principal and Guest Investigators utilized the spectrograph to observe some 105 astronomical targets. These data will enter the public domain in early to mid 1998.  Lyman/FUSE is not far from launch, and will offer a much higher spectral resolution and sensitivity for point sources. However, near-trivial modifications would enable the Berkeley spectrograph to achieve a significantly superior performance for studies of extended emission. New replicas of the current gratings and SiC overcoating of the primary mirror would allow the spectrograph to achieve 2 \\AA\\ slit-limited resolution through a 45 $\\times$ 420\\arcsec\\ slit and an effective area of $\\sim$ 50 cm$^2$ across the 900 -- 1200 \\AA\\ band. The angular resolution along the slit length would be better than 1 \\arcmin. Such an instrument would be highly desirable for studies of intracluster gas, galaxies and their halos, supernova remnants, and other extended objects. At time of writing, there are no specific plans for a third flight of the payload."
    },
    "9804/astro-ph9804283_arXiv.txt": {
        "abstract": "In this paper we propose a new statistic capable of detecting non-Gaussianity in the CMB. The statistic is defined in Fourier space, and therefore naturally separates angular scales. It consists of taking another Fourier transform, in angle, over the Fourier modes within a given ring of scales. Like other Fourier space statistics, our statistic outdoes more conventional methods when faced with combinations of Gaussian processes (be they noise or signal) and a non-Gaussian signal which dominates only on some scales. However, unlike previous efforts along these lines, our statistic is successful in recognizing multiple non-Gaussian patterns in a single field. We discuss various applications, in which the Gaussian component may be noise or primordial signal, and the non-Gaussian component may be a cosmic string map, or some geometrical construction mimicking, say, small scale dust maps. ",
        "introduction": "Current theories of structure formation may be roughly divided into two classes: active and passive perturbations. According to the inflationary paradigm, quantum fluctuations in the very early universe are produced during a period of inflation \\cite{cosmcross} and grow to become classical density perturbations \\cite{lidlyth,bardeen}. These perturbations evolve linearly until late times when the overdensities become galaxies. Perturbations due to inflation are called passive because they are seeded at some initial time near the Planck time and then evolve `deterministically', or linearly. They leave their imprint on the CMB at last scattering \\cite{pyu,nature,HS1,HS2,richard,1.3}. Although this is not strictly necessary, in most cases  the fluctuations in the CMB temperature due to inflationary perturbations form a Gaussian random field \\cite{bondefst,bardstat}.  There is another class of theories of structure formation, topological defects caused by phase transitions in the early universe \\cite{Kib,vv}. Perturbations caused by defects are known as active perturbations \\cite{mafc,coher} since they are continually being seeded by an evolving network of defects through the history of the universe. The fluctuations in the CMB in defect models have been found to be non-Gaussian \\cite{phases}, even though the extent and strength of this non-Gaussianity is still far from clear.  Recently a wide class of defect models have been shown to be in conflict with current data \\cite{neil,abr,abr2} but some viable active models still remain \\cite{mimic,durrer,abrw}. In fact, some of the most interesting \\cite{abr3,shellard} of these (based on cosmic strings) require a non-zero cosmological constant of the sort that currently favoured by supernova experiments \\cite{sn1}.   Thus it is important to be able to distinguish Gaussian from \\ngn fluctuations. Many different tests are being tried \\cite{Kogut,XLuo,fermag,cumul,gorski,sergei}, adapted to different experimental settings, and types of signal. Our statistic is designed to be used for small fields where one or more distinctive shapes are obscured by an extra Gaussian component, and so are not visible in real space. It is our experience in previous work \\cite{fermag} that in such situations standard statistics, based in real space, fail to recognize the non-Gaussianity of the signal. The idea is to study the statistical properties of map derivatives which are only sensitive to a given scale. In this way we may separate out different scales, some of which may have Gaussian or nearly Gaussian fluctuations, some very non-Gaussian. Although this scale filtering may be achieved using the wavelet transform \\cite{cumul,pando}, in this paper we choose to use the Fourier transform as our scale filter.  One problem with this approach is that the Fourier transform is a global transformation, and therefore can only recognize non-Gaussian structures globally. It may therefore offer a rather contorted description for complicated networks made up of essentially simple objects. We will however perform a transformation over the Fourier modes: a second Fourier transform, in angle, for modes in each ring in Fourier space. We will argue that by doing so we are able to recognize mostly features of individual objects, and so bypass this problem. This operation returns a set of quantities which are blind to the random orientations of individual objects. Unfortunately their random positions still affect our statistic, so there will be a limit (albeit less stringent) on the number of objects in the field before our statistic becomes confused.      The plan of this paper is as follows. In Section II we look at some of the ideas involved in thinking about non-Gaussianity. In Section III we define our statistic and look at some of its properties and motivations. In Section IV we apply our statistic to practical situations in which subtle non-Gaussian signals are present. We consider non-Gaussian signals corrupted by the presence of a Gaussian process, which can be noise or primordial signal, and which dominates on all but a narrow band of scales. We consider two types of non-Gaussian signal. We consider CMB maps obtained from string simulations, implementing the algorithms in \\cite{av}. We also consider geometrical constructions mimicking, say, small scale dust maps. We show how such maps look very Gaussian, but fail to confuse our statistic. ",
        "conclusions": "We have introduced a statistic to look for \\ng in the cosmic microwave background. It is designed to pick out \\ngn features which are superimposed on Gaussian fluctuations and which are therefore not visible in real space. Since some scales may be more \\ngn than others we choose to separate them out. It is our experience in previous work \\cite{fermag} that in such situations standard statistics, based in real space, fail to recognize the non-Gaussianity of the signal.  Our statistic is naturally tailored for interferometric experiments, which make measurements in Fourier space. Indeed the algorithm exposed above could easily be turned into a data analysis package operating over visibilities.  One of the problems with Fourier space statistics is that they recognize only global shapes. Therefore they become very ineffective when many individual structures are present. By taking another Fourier transform, in angle, over modes in a given ring in Fourier space, we can factor out orientations, and be only sensitive to an average shape of individual structures. Their random positions, on the other hand, will eventually make our statistic very ineffective as the number of objects becomes large. We found that although we have improved on previous work, still there is a limit on how many individual structures there may be in the field before the statistic fails to pick out their non-Gaussianity.  Finally, some comments are in order regarding the applications we have considered. The Gaussian component we have considered may include noise, if it has the right power spectrum. Hence we have already considered the effects of a simplified form of noise: it will merely reduce the band of scales where the non-Gaussian signal dominates, typically providing it with an upper boundary. An exception is the case of non uniform noise in the u-v plane, present in most interferometer experiments. Non uniform noise in the Fourier domain will look non-Gaussian, for it is a Gaussian process which is not isotropic or translationally invariant. Therefore an extra element of confusion, not dealt with in this paper, will appear.  Regarding the non-Gaussian component we have considered two types of signal. We have looked at some of the properties of the statistic when it is applied to a field containing one or more distinct shapes. The geometrical constructions presented are somewhat reminiscent of the jagged structures present in small scale dust maps. Experimenting on fields containing many structures we saw our statistic rapidly become blind to their non-Gaussianity. However interferometer fields are often small. They would therefore contain only a few, but more than one, of these structures. This is something our statistic can cope with, unlike previous Fourier space based statistics.  We have also studied our statistic when applied to cosmic string maps. We found that some individual small fields showed \\ng in the signal. If the results were averaged over many fields the effect of \\ng on $|F(\\beta)|$ would not show. Once again this strategy is ideal for interferometers, for which combining small fields into a large field is in fact a rather awkward operation."
    },
    "9804/astro-ph9804230_arXiv.txt": {
        "abstract": " ",
        "introduction": "La d\\'ecouverte en 1965 par Penzias et Wilson du fond de rayonnement cosmologique  diffus \\`a 3 K a d\\'efinitivement fait du mod\\`ele du Big Bang chaud le cadre g\\'en\\'eral de la cosmologie moderne.  Ces trois derni\\`eres d\\'ecennies ont \\'et\\'e marqu\\'ees par de  nombreux progr\\`es tant th\\'eoriques qu'observationnels qui n'ont fait que confirmer sa validit\\'e, et de pr\\'eciser petit \\`a petit les d\\'etails de ce sc\\'enario.  Parmi les avanc\\'ees notables de ces derni\\`eres ann\\'ees, citons pour m\\'emoire:  \\begin{itemize} \\item La r\\'ealisation de grands catalogues de galaxies. On a maintenant des catalogues tri-dimensionnels qui contiennent des milliers d'objets. Ces catalogues permettent de faire une v\\'eritable cosmographie de l'univers local. De nouveaux moyens sont apparus r\\'ecemment pour compl\\'eter ce type de catalogues, la d\\'etermination des champs de vitesse cosmiques, permet ainsi d'acc\\'eder \\`a des informations cin\\'ematiques tr\\`es pr\\'ecieuses. Et depuis la fin des ann\\'ees 80, un nouveau moyen d'investigation est en train d'\\'emerger, il s'agit des cartes de distorsion gravitationnelle. Elles permettent de visualiser les lignes de potentiel de la masse projet\\'ee.  \\item le mod\\`ele de Mati\\`ere Noire Froide (CDM pour Cold Dark Matter) qui est apparu au d\\'ebut des ann\\'ees 80 et qui a servi de point de r\\'ef\\'erence (\\`a d\\'efaut de devenir un mod\\`ele standard) pour tous les travaux sur le probl\\`eme de la formation des grandes structures.  \\item Le d\\'eveloppement des th\\'eories inflationnaires. C'est encore un terrain tr\\`es sp\\'eculatif, mais c'est le lieu tr\\`es excitant o\\`u les concepts de la physique des hautes \\'energies rencontrent des exigences observationnelles de plus en plus fiables.  \\item Enfin, la d\\'etection en 1992 des fluctuations de temp\\'erature du fond de rayonnement cosmologique par l'exp\\'erience satellitaire COBE/DMR (Smoot et al. 1992) a marqu\\'e un tournant pour la cosmologie: pour la premi\\`ere fois on avait une preuve directe de l'origine des grandes structures de l'univers.  \\end{itemize}   Dans ce cours je vais principalement m'int\\'eresser au probl\\`eme de la formation des grandes structures.  Quelques ouvrages de r\\'ef\\'erence, \\begin{itemize} \\item Relativit\\'e G\\'en\\'erale: Weinberg, {\\it Gravitation and Cosmology}, 1972; Landau et Lifschitz, {\\it Classical Theory of Fields}, 1975 \\item Univers primordial: Kolb et Turner {\\it The Early Universe}, 1990, mais il est peu pr\\'ecis; revue de Brandanberger, {\\it Inflation and Cosmic Strings: two Mechanisms for producing Structure in the Universe}, Int. J. Mod. Phys. A2 : 77, 1987. \\item Inflation: Linde, {\\it Particle Physics and Inflationary Cosmology}, 1990; Liddle et Lyth, {\\it The Cold Dark Matter density perturbation}, 1993, Physics Reports, 231, 1; Lidsey et al. {\\it Reconstructing the inflaton potential - an overview}, 1997, Reviews of Modern Physics, 69, 2 \\item Formation des grandes structures: Peebles {\\it The Large Scale Structure of the Universe}, 1980; {\\it Principle of Physical Cosmology}, 1993; V. Sahni \\& P. Coles, {\\it Approximation Methods for Non-linear Gravitational Clustering}, 1995, Physics Reports, 262, 1 \\end{itemize} ",
        "conclusions": "Dans ce petit tour d'horizon de la cosmologie j'ai \\'et\\'e loin d'\\^etre exhaustif. Dans le domaine de l'\\'evolution des grandes structures je n'ai pas mentionn\\'e bon nombre de domaines de recherche en plein d\\'eveloppement, comme \\begin{itemize} \\item L'\\'etude des flots cosmiques \\`a grande \\'echelle, et leur utilisation pour mesurer les param\\`etres cosmologiques; \\item L'\\'etude des amas de galaxies et de leur contenu aussi bien en mati\\`ere noire (\\'etudes dynamique, reconstruction de masse par effets de lentilles gravitationnelles..), qu'en mati\\`ere baryonique (galaxies, rayonnement X..); \\item L'\\'etude des absorbants et des objets de grand $z$; \\item ... \\end{itemize}  Du point de vue du physicien th\\'eoricien un domaine de recherche de pr\\'edilection est inconstablement la cosmologie primordiale, lieu de rencontre entre entre la physique des hautes \\'energies et la cosmologie observationnelle.  Cependant la formation des grandes structures recelle un certain nombre de probl\\`emes ouverts digne d'int\\'er\\^et, \\begin{itemize} \\item La th\\'eorie des perturbations est loin d'avoir livr\\'e tous ses secrets. En particulier il serait extremement utile de comprendre les corrections en boucles. \\item De nombreux aspects du r\\'egime non-lin\\'eaire ne sont pas compris. Par exemple on ne sait pas d\\'ecrire la transition vers le r\\'egime multiflots qui conduit \\`a la virialisation de la mati\\`ere dans les puits de potentiel. Il serait aussi tr\\`es int\\'eressant de pouvoir exhiber une solution explicite des \\'equations dynamiques dans le r\\'egime nonlin\\'eaire (m\\^eme si ce n'est qu'une forme assymptotique). \\item L'exploitation des catalogues tridimensionnel ou bidimensionnel notamment avec la mise en \\'evidence de propri\\'et\\'es non-Gaussiennes n'est pas encore optimale. La question se pose par exemple pour les cartes de distorsion gravitationnelle. \\item La relation entre baryons et mati\\`ere noire est un domaine qui n'a pratiquement pas \\'et\\'e explor\\'e. On ne dispose pour l'instant que d'exp\\'eriences num\\'eriques pour essayer de comprendre ce qui se passe. \\end{itemize}"
    },
    "9804/astro-ph9804006_arXiv.txt": {
        "abstract": "The Lyman edge at 912 {\\AA} is an important diagnostic region for studying quasi-stellar objects (QSOs).  In particular, it reveals a great deal about the physical conditions within the atmospheres of accretion disks, a ubiquitous component of QSO theories.  A robust prediction of accretion disk models is a significant polarization due to electron scattering just longward (in wavelength) of the Lyman edge because of the wavelength dependence of the Hydrogen absorption opacity.  Observations of the Lyman edge regions of QSOs have shown scant evidence for the predicted features---few QSOs show the broad, partial Lyman edges expected to be common according to most theories, and none show the high polarizations expected longward of the Lyman edge.  Still, polarization spectra of a small number of QSOs have shown a rising polarization (up to 20\\%) at wavelengths {\\it shortward} of the Lyman edge. We have now doubled our sample of intermediate-redshift QSOs observed with the {\\it HST/FOS} spectropolarimeter to determine the amount of polarization on both sides of the Lyman limit. For this new sample of six objects, polarizations are low and mostly consistent with zero below the Lyman edge.  Another important result of the new data is that it strengthens the conclusion that quasars are generally not polarized significantly just longward of the Lyman edge at $\\sim$1000\\AA\\/.  There is no significant statistical wavelength dependence to the polarization longward of the Lyman edge indicating that simple plane-parallel atmospheres with scattering-dominated opacity are not significant sources of UV flux in quasars. ",
        "introduction": "One of the fundamental components of most theories of quasi-stellar objects (QSOs) is an accretion disk.  As gas is fed into the central regions of the QSO, residual angular momentum causes the gas to naturally settle into a disk.  While most theories predict the formation of such a disk, few address the observational consequences of the disk models in detail.  These theoretical studies have found that the Lyman edge at 912 {\\AA} is a powerful diagnostic feature for the physical characteristics of the disk.  The simplest disk models (quasi-static with viscous dissipation at large optical depth) predict Lyman edges in either emission or absorption, depending on the viewing angle and the physical details of the disk atmosphere.  Such Lyman edges would be broadened by rotation of the disk and by general relativistic effects as light passes close to the central black hole.  In most QSOs such edges are not seen, although Koratkar, Kinney, \\& Bohlin (1992) found a small number of candidate ``partial edges'' in IUE data.  A second disk signature is the linear polarization, $P$, of the continuum from the disk.  A purely scattering atmosphere will produce high polarization perpendicular to the disk axis.  Again, this signature is not seen in any QSOs; in fact generally QSOs show {\\it low} optical polarization {\\it parallel} to the inferred disk axis. Laor, Netzer, \\& Piran (1990) attempted to show that a disk atmosphere should have significant absorptive opacity, and thus can produce dramatically lower optical polarization (albeit still perpendicular to the disk axis and thus inconsistent with the observations).  A more robust prediction according to their work, however, is a rise in polarization with decreasing wavelength from the optical into the UV. Just longward of the Lyman edge $P$ is highest, since it is at these wavelengths that scattering best competes with absorption.  Just shortward of the Lyman edge, as the absorption opacity increases, $P$ should drop again.  According to Laor \\etal, this polarization signature should appear even when no disk signature is seen in total flux.  In our first polarization study of three of the rare objects known from IUE spectra to have partial Lyman edges at the systemic redshifts, we found low polarizations longward of the edges, so we did not confirm the Laor \\ea prediction.  We surprisingly did find high polarization {\\it shortward} of the edge in a couple of objects, contrary to the accretion disk predictions of Laor \\etal, but qualitatively explicable by effects found in the more detailed calculations of Blaes \\& Agol (1996);  see also Agol \\& Blaes (1996).  The previous studies of intermediate redshift quasars (Koratkar \\ea 1995 and Impey \\ea 1995), included 4 objects, one of which shows only a marginal detection of polarization, while the remaining three objects show significant polarization ($>$ a few percent) shortward of the Lyman edge. A study of three high redshift objects observed from the ground failed to show any polarization changes at the edge position; most had tight limits on the  polarization longward of the edge, but noisy data shortward of the Lyman edge (Antonucci \\ea 1996).  PG 1630+377 is the only object yet observed that can be studied in any detail (Koratkar \\ea 1995; Paper I).  In this object $P$ rises rapidly shortward of the edge, reaching 20\\% by 1600 {\\AA} (650 {\\AA} rest wavelength).  The Ly$\\alpha$ emission line also shows a high (7.3\\%) polarization at the same position angle. Antonucci \\ea (1996) discuss polarization observations and other constraints on disk models in some detail.  Based on the small number of QSOs observed in the UV in polarization, at the time of paper I, we could say little about whether high polarization shortward of the Lyman edge is common.  Hence, we have significantly expanded the UV polarization database by observing six bright, $z > 1$ QSOs both below and above the Lyman edge. In this paper we discuss these new spectropolarimetric ultraviolet observations from the {\\it Hubble Space Telescope} Faint Object Spectrograph ({\\it HST/FOS}).  A difference with respect to our previous study, however, is that only two of the new objects were suspected to have partial edges in total flux at the systemic redshift.  The rest were simply selected because they show significant flux at short wavelengths. ",
        "conclusions": "Of the six QSOs identified by Koratkar, Kinney \\& Bohlin (1992) as candidate targets which have partial Lyman edges consistent with edges from simple thin accretion disks, we now have spectropolarimetric observations for five QSOs. We showed in section 3.2 that one of these candidates, 0743$-$673 no longer qualifies as a partial Lyman edge object.  There are only 13 high and intermediate redshift QSOs which have spectropolarimetry observations shortward of the Lyman edge region. These objects come from this paper, Koratkar \\ea (1995), Impey \\ea (1995), and Antonucci \\ea (1996).  At this point any detailed statistical tests of polarization distributions are certainly not warranted given the inhomogeneous selection criteria and data quality, and the highly model-dependent predictions. Yet, Lyman edge spectropolarimetry results can be summarized as follows:  \\begin{itemize}  \\item{}Of the 13 objects only three objects (PG 1630+377, PG 1338+416 and PG 1222+228 from paper I and Impey \\ea 1995) show significant polarization at wavelengths shorter than 912\\AA\\/ (Lyman edge).  To these three objects we can add one more marginal detection (PKS 0405$-$123 from paper I). All 13 objects in the sample show a polarization signature which is inconsistent with any simple accretion disk model. Furthermore, $\\sim$30\\% of the sample show a rise in polarization shortward of the Lyman edge.  This observed rise in polarization is qualitatively consistent with the disk models of Blaes \\& Agol (1996).  A number of different interpretations of the UV signature have been given by Lee \\& Blandford (1997), Shields, Wobus \\& Husfeld (1997) and by us in paper I.  We urge the interested reader to consult those papers for more details.  \\item{}There are a total of five objects (PG 0117+213 from the present sample, PG1630+377, PG 1338+416, and PKS 0405$-$123 from paper I, and 0014+813 from Antonucci \\ea 1996) which show candidate partial Lyman absorption edges due to accretion disks at the systemic redshift in total flux.  Of these five objects, two (PG 1630+377, PG 1338+416) have sufficient signal-to-noise at rest wavelengths of $\\leq$750\\AA\\/, and show significant polarization shortward of the Lyman edge (at least a few percent, detected at four sigma or greater significance). If the rise in UV polarization seen in PG 1630+377 is characteristic of objects with partial Lyman edges we need to observe rest wavelengths as short as $\\sim$700\\AA.  The effective shortest rest wavelength observed in PG 0117+213 and PKS 0405$-$123 is $\\sim$800\\AA\\/. Thus in these objects we could have missed the rise in polarization, although we do have a marginal detection for PKS 0405$-$123. 0014+813 does not have sufficient signal-to-noise shortward of the Lyman edge. To summarize, of the objects that show candidate partial Lyman absorption edges in total flux, $\\sim$40\\% show polarization shortward of the Lyman edge.  \\item{}We have eight objects in the sample of 13 that do not show a partial Lyman edge feature in total flux. Only one object out of these, PG 1222+228, from Impey \\ea (1995) shows significant polarization shortward of the Lyman edge ($P$ = 4.6$\\pm$0.9\\%). The rest show no detection of polarization in the 912\\AA\\/ spectral region. The linear polarization upper limits shortward of the Lyman edge in the remaining objects is \\ltsima 4\\%.  \\item{}PG 1630+377 is the best studied object (see Paper I), and shows UV polarization reaching 20\\% at 650\\AA\\/ rest wavelength. Such high degree of polarization is rare in (non-blazar) QSOs.  \\end{itemize}  Perhaps the simplest result of the current data is that it strengthens the conclusion that quasars are generally not polarized significantly just longward of the Lyman edge at $\\sim$1000\\AA\\/.  This paper, Koratkar \\ea (1995), Impey \\ea (1995), and Antonucci \\ea (1996) together present good observations of about 20 objects in the region just longward of the Lyman edge, all with low UV polarization (\\ltsima 1.5\\%).  Further, there is no significant statistical wavelength dependence to the polarization as predicted by the models of Laor et al (1990). From these observations we conclude that simple plane-parallel atmospheres with scattering-dominated opacity are not significant sources of UV flux in quasars.  Recapitulating the previous discussions here briefly, we note that models from the 1980s generally assumed that AGN accretion disks are powered by viscous dissipation below the atmospheres, and that the atmospheric opacities are dominated by electron scattering opacity in the annuli that produce the rest optical and UV.  This results in 0\\% to 11.7\\% polarization (Chandrasekhar 1960), depending on inclination, and in a direction perpendicular to the symmetry axis of the disks. Pioneering optical polarimetry observations showed much smaller polarizations, which are {\\it parallel} to the axes when the latter could be inferred from a radio jet position angle (Stockman, Angel \\& Miley 1979; Antonucci 1988).  To explain the low observed polarization, subsequent models by Laor et al (1990) suggested that electron scattering was only important in the $\\sim$1000-2000 \\AA\\/ range, with the Lyman continuum and free-free absorption opacity dominating shortwards and longwards of that interval respectively.  Other more detailed calculations revealed that a lower fraction of absorption opacity was required to reduce the predicted polarization than was assumed by Laor et al; and that under rather special circumstances a large polarization {\\it parallel} to the disk axis could be produced shortward of the Lyman edge (Blaes and Agol 1996, and references therein).  The observed rise in UV polarization shortward of the Lyman edge has been interpreted both in the context of accretion disk models and non-disk related models. Here we do not further discuss the polarization and depolarization mechanisms discussed in detail in Paper I. An additional key complication in the accretion disk models, is that AGN variability data require that the disk atmosphere is actually illuminated from above, (e.g. Antonucci 1988, Sincell and Krolik 1996), perhaps producing a strong polarization which cannot be calculated rigorously without specification of the illumation geometry."
    },
    "9804/gr-qc9804041_arXiv.txt": {
        "abstract": "{The mass function of primordial black holes created through the near-critical gravitational collapse is calculated in a manner fairly independent of the statistical distribution of underlying density fluctuation, assuming that it has a sharp peak on a specific scale.  Comparing it with various cosmological constraints on their mass spectrum, some newly excluded range  is found in the volume fraction of the region collapsing into black holes as a function of the horizon mass. }  \\pacs{PACS Numbers: 04.70.Bw, 04.70Dy, 98.80.Cq } ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804140_arXiv.txt": {
        "abstract": "We investigate whether models based on the assumption that jets in quasars are powered by rotating black holes can explain the observed radio dichotomy of quasars. We show that in terms of the ``spin paradigm'' models, radio-loud quasars could be objects in which the black hole's rotation rate corresponds to an equilibrium between spin-up by accretion and spin-down by the Blandford-Znajek mechanism. Radio-quiet quasars could be hosting black holes with an average spin much smaller than the equilibrium one. We discuss possible accretion scenarios which can lead to such a bimodal distribution of black hole spins. ",
        "introduction": "Quasars are characterized not only by an intense radiation from the central engine, but also by jets which power large scale radio structures. The ratio of radio luminosity of these structures to the optical luminosity of the central sources shows a bimodal distribution, with only $\\sim 10$ \\% of quasars belonging to the radio-loud category (see, e.g., Kellerman et al. 1989; Hooper et al. 1995; Falcke, Sherwood, \\& Patnaik 1996; Bischof \\& Becker 1997). Radio-loud quasars have never been found in spiral galaxies, while the hosts of radio-quiet quasars are either spiral or elliptical galaxies (Taylor et al. 1996; Kukula et al. 1998).  On the other hand, both radio-quiet and radio-loud quasars have almost identical average IR-optical-UV spectra (Francis et al. 1993; Zheng et al. 1997), which suggests similar accretion conditions in these two samples of objects.  These properties seem to support models which are based on the idea that jets are powered by rotating black holes.  However, in order to explain the observed radio loudness bimodality in terms of a bimodal distribution of black hole spins, one has to assume that the population of supermassive black holes is dominated by very low-spin black holes, and one should understand why radio-loud quasars are always hosted by elliptical galaxies.  This issue has been addressed by Wilson \\& Colbert \\shortcite{wc95}, who proposed that high-spin black holes exist only if formed by coalescence of two black holes. Such coalescences would take place mostly in a denser galactic environment (groups, clusters), which are much more populated by elliptical galaxies than the field regions.  However, the basic assumption of this scenario, that black holes which do not undergo coalescence rotate very slowly -- despite the angular momentum gained from accretion discs -- must be verified. In other words, one should verify that there exists a mechanism which could keep black holes at low rotation rates despite of the accreted angular momentum.  This problem was investigated by Moderski \\& Sikora \\shortcite{ms96a}. They showed that the Blandford-Znajek (B-Z) mechanism, which extracts rotational energy from a black hole \\cite{bz77}, is not efficient enough to counteract the spinning up of a black hole by the gas accreted from standard $\\alpha$-discs \\cite{ss73}.  Low-spin equilibrium states are possible only for very low accretion rates but the time required to established such an equilibrium is longer than the age of the Universe.  This is because the power extracted from the black hole is proportional to the square of intensity of the magnetic field which threads the black hole.  The field intensity is limited by the pressure in the accretion flow and, therefore, is very low for very low accretion rates. Hence, once a black hole gets a high spin, it will be rotating fast forever.  Low-spin equilibrium solutions for high accretion rates are possible only for the so called $\\beta$-discs \\cite{ms97}.  In such discs the viscous stress driving accretion is proportional to the gas pressure only (and not to the total pressure as in the case of $\\alpha$ discs). Since the relation between the viscous stress and the pressure in radiation pressure dominated accretion discs is unknown, such discs are a reasonable alternative to $\\alpha$-discs. For high accretion rates the gas pressure is by several orders of magnitude smaller than the radiation pressure \\cite{ss73}.  This implies that, for a given accretion rate, $\\beta$-discs are much denser than $\\alpha$-discs \\cite{sc81} and, therefore, can confine much stronger black hole magnetic fields.  Since the power extracted from a black hole via the B-Z mechanism scales with the square of the magnetic field intensity, equilibrium spins for $\\beta$-discs are smaller than for $\\alpha$-discs. This, of course, does not mean that the power extracted from a black hole with a $\\beta$-disc must be smaller than the power extracted from a black hole with an $\\alpha$-disc, because in the expression for the B-Z power a lower equilibrium spin is compensated by a higher intensity of the magnetic field \\cite{ms97}. Thus, low equilibrium spin black holes can represent radio-quiet quasars only, if the conversion of the extracted black hole energy into jet energy is very inefficient. According to the Wilson-Colbert scenario powerful jets would then exist only in objects where a coalescence of two black holes leads to spins much higher than the equilibrium one. However, powers extracted from black holes with spins much larger than the equilibrium one are so large, that such black holes would be spun-down on time scales two orders of magnitude shorter than the typical lifetime, $\\sim 10^8$ years, of radio quasars (see, e.g., Leahy, Muxlow \\& Stephens 1989).  Powering jets in radio-loud quasars by a black hole can last as long as $10^8$ years, only if losses of angular momentum due to the B-Z mechanism are compensated by gains of angular momentum from an accretion disc. This would give a correlation between radio and optical luminosities, in accordance with observations \\cite{ser97}. Then, however, one would have to explain why the majority of super-massive black holes have spin values much lower than the equilibrium one -- the condition for the existence of radio-quiet quasars. As suggested by Moderski, Sikora \\& Lasota \\shortcite{msl97}, black holes in most objects could be forced to rotate slowly by multi-accretion events with random orientations of the angular momentum vector. In such a scenario, quasars which become radio-loud are only those which undergo major accretion events, induced, i.e., by a merger of two big galaxies.  Following this, a black hole can easily double its mass and reach an equilibrium state.  In this paper we explore this possibility and derive conditions the model must satisfy in order to explain the observed radio properties of quasars.  This paper is organized as follows. In Section 2 we present equilibrium spin solutions, obtained for a variety of accretion disc models. In Section 3, we discuss multi-event accretion scenarios. In Section 4 we use our results to derive conditions which must be satisfied by quasar evolution models in order to obtain bimodal distribution of black hole spins. ",
        "conclusions": "Radio dichotomy of quasars was discovered many years ago (Strittmatter et al. 1980; Kellerman et al. 1989), but is still waiting for a theoretical explanation.  As for now, the consensus concerns only one aspect of the problem: it is clear that jets in quasars must be formed near a supermassive black hole.  This follows from the energetics of quasar jets, since no other known sources could power jets at a rate reaching $10^{46}$ ergs s$^{-1}$ for millions of years (Rawlings \\& Saunders 1991; Leahy et al. 1989). Independent argument for the formation of extra galactic jets in the vicinity of supermassive black holes is provided by direct VLBI observations of nearby radio galaxies.  In particular, in 3C 274 (M87) the jet is seen down to $10^{16}$ cm from the center \\cite{jb95}, which corresponds to $100$ gravitational radii for the $3 \\times 10^9 M_{\\odot}$ \\cite{har94} central black hole.  There, very deep in the gravitational potential well, jets could be powered either by the innermost parts of an accretion disc (Blandford \\& Payne 1982; Park \\& Vishniac 1994; Contopoulos 1995; Begelman 1995) or by a rotating black hole (Blandford \\& Znajek 1977; Rees et al. 1982).  However, no jet production model can be successful, if it fails to explain why only a small fraction of quasars is radio loud and why radio-loudness has a bimodal distribution.  In terms of the $R=F_r/F_o$ ratio, where $F_r$ and $F_o$ are the monochromatic fluxes measured at frequencies $\\sim 10^{10}$ Hz and $\\sim 10^{15}$ Hz, respectively, radio-quiet quasars cluster around $R \\sim 0.3$ and radio-loud quasars cluster around $R \\sim 300$ (Kellermann et al. 1989; Falcke, Sherwood \\& Patnaik 1996).  Thus, the average radio-loudness of the two quasar populations differs by a factor $10^3$ and this number, together with typical radio luminosities of radio loud quasars, $L_r \\sim 10^{45}$ ergs s$^{-1}$, provides the basic quantitative conditions which should be satisfied by any unified model of quasars.  These conditions, together with our results discussed in the two previous sections are used below to test spin based models of a jet activity in quasars. The predictions of such models should also satisfy such observationally established trends, as  \\noindent - radio-loud quasars avoid  disc-galaxies and have UV-luminosities $ \\ge 10^{46}$ ergs s$^{-1}$;  \\noindent - radio-quiet quasars are present both in spiral and elliptical galaxies \\cite{tay96} and their radio properties do not depend on the galaxy morphology \\cite{kuk98};  \\noindent - radio properties of radio-quiet quasars suggest that they are, like in radio loud quasars,  related to the jet production by a central engine.  Assuming that the efficiency of conversion of jet energy into radio emission is $\\sim 10$\\%, the typical jet in radio-loud quasars should have $P~\\sim~10^{46}$ ergs s$^{-1}$. Similar jet powers are deduced by calculating the total energy content of extended radio sources and dividing it by the age of the source \\cite{lea89}, or from energetics of $\\gamma$-ray production in sub-parsec jets (see, e.g., Sikora 1997).  Largest powers which can be extracted from rotating black holes are given by equation (\\ref{power}). For $A=1$ and $B_{\\perp} = 8\\pi p_{tot}$ we obtain $P_{max} \\sim 3 \\times 10^{44} M_8^2 p_{tot,8}$ ergs s$^{-1}$, where $p_{tot,8} = p_{tot}/10^8$dyne cm$^{-2}$ and $M_8 = M/10^8 M_{\\odot}$.  One can see from Figure~\\ref{pmax} that for high accretion rates ($\\dot m > 0.1$, say) and black hole masses $\\sim M=10^9 M_{\\odot}$, a pressure $\\ge 10^8$ dyne cm$^{-2}$ is provided by $\\alpha$-discs with $\\alpha \\le 0.1$, and by all $\\beta$-discs.  Thus, the B-Z mechanism is efficient enough to power jets in radio-loud quasars, provided the black hole magnetic field is supported by the total disc pressure.  If the latter is not true and, as Ghosh and Abramowicz \\shortcite{ga97} argued, the energy density of the black hole magnetic field cannot exceed the energy density of the maximum magnetic field in a disc, then the maximum pressure of the black hole's magnetic field is numerically equal to the total pressure in an $\\alpha=1$ disc. In this case black hole masses $\\sim 3 \\times 10^9 M_{\\odot}$ are required in order to get $P \\sim 10^{46}$ ergs s$^{-1}$. The case of M87 seems to prove that such black holes are not necessarily exceptional \\cite{har94}.  However, one should note here, that the question of the diffusion of an external magnetic field into an accretion disc is still open (see, e.g., Wang 1995; Bardou \\& Heyvaerts 1996).  Assuming, as before, that the fraction of the jet energy converted into radiation is 10\\%, and that the bolometric corrections for jet radiation at $\\sim 10^{10}$Hz and for accretion disc radiation at $\\sim 10^{15}$Hz are of the same order, we obtain that $P/L_d \\sim 10 L_r/L_o \\sim 10 F_r{\\nu}_r/F_o{\\nu}_o \\sim 10^{-4} R$. Thus, for radio-loud quasars models should predict $P/L_d \\sim 0.1$, while radio-quiet quasars should cluster around $P/L_d \\sim 10^{-4}$. As is seen from Figures~2 and 3, $P/L_d \\sim 0.1$ nicely corresponds to black hole equilibrium spin solutions for all but $\\alpha > 0.1$ disc models. One can also check, that there are no equilibrium spin solutions which would correspond to radio-loudness of radio-quiet objects. For them $A < 0.03$ is required, provided that radio luminosity scales linearly with $P$. A population of such low spin black holes can exist only if black holes are born with very low spin and then accrete very little (Moderski, Sikora \\& Lasota 1997), or if black hole evolution is determined by multi-accretion events with random angular momenta.  As one can deduce from results presented in Figure~7, hundreds of accretion events per object are required in order to have more than $90$ \\% of black holes with $A < 0.1$ at any given moment. This is too much to be obtained by accretion events induced by capture of dwarf galaxies, but can be achieved by accretion of molecular clouds. Molecular cloud accretion events were recently proposed by Sanders \\shortcite{san98} to explain some properties of Sgr A$^*$ and other AGNs.  This scenario is supported by the random orientation of central engines vs. the orientation of galactic discs, as deduced from observations of ``UV'' cones (Wilson \\& Tsvetanov 1994; McLeod \\& Rieke 1995) and radio axis \\cite{ckp98} in Seyfert galaxies. Here we should note, that in our simplified treatment of the multi-accretion scenario (Section~3), we didn't take into account the coupling between the spin of the black hole and the orbital angular momentum of the approaching molecular clouds. Such coupling supposedly leads to random wondering of the black hole spin vector.  One can now speculate that changes of orientation of the black hole spin could be interrupted and the black hole could be spun-up to very high spins following a merger process. This process could induce a massive and long lasting accretion event. If during such an event the accretion proceeds from a fixed plane and at least doubles the black hole mass, the black hole spin reaches the equilibrium spin and the object becomes a typical radio-loud quasar \\cite{msl97}.  Since mergers happen mostly in groups and clusters of galaxies, where the population of galaxies is dominated by ellipticals, this could explain why radio loud quasars avoid spiral galaxies. Observational arguments for such a scenario are exactly the same as those used by Wilson and Colbert \\shortcite{wc95}. The only difference is, that they postulated formation of high spin black holes via coalescence of two supermassive black holes, while in our scenario high spins result from an accretion process. Note, however, that a coalescence of two black holes, if it happens, does not have to affect much our scenario. If the coalescence involves two black with very different masses, the final spin will be determined by the accretion process, otherwise both processes lead to similar spins.  What are the perspectives for an observational test of the assumption that radio-quiet objects have low spins?  A possibility to measure the spin of supermassive black holes is provided by the detailed studies of profiles of the X-ray fluorescent iron line produced in the surface layer of the innermost parts of accretion discs. Such lines are detected in many Seyfert galaxies, which represent the low luminosity branch of radio-quiet quasars.  For at least one of such objects the line profile was claimed to be consistent with the kinematics given by the rotation of a disc around a black hole in fast rotation \\cite{iwa96}. However, as demonstrated by Reynolds and Begelman \\shortcite{rb97}, similar line profiles can be produced around non-rotating black holes, provided that a large part of the line emission comes from below the marginally stable orbit.  Therefore, much more detailed theoretical models and sensitive observations are required to get conclusive diagnostics from this type of investigations.  The remarkable discovery of relativistic jets in several Galactic X-ray sources (cf. Mirabel \\& Rodriguez 1994; Hjellming \\& Rupen 1995; Newell, Spencer \\& Garrett 1997) suggests that the radio-dichotomy exists for Galactic compact objects as well. As was argued recently by Zhang, Cui \\& Chen \\shortcite{zcc97}, jet activity in these sources can also be conditioned by the value of the black hole spin."
    },
    "9804/astro-ph9804154_arXiv.txt": {
        "abstract": "We study the interpretation of the mean surface density of stellar companions as a function of separation (or, equivalently, the two point correlation function of stars) in star-forming regions. First, we consider the form of the functions for various simple stellar distributions (binaries, global density profiles, clusters, and fractals) and the effects of survey boundaries.  Following this, we study the dependencies of the separation at which a transition from the binary to the large-scale clustering regime occurs.  Larson \\shortcite{Larson95} found that the mean surface density of companions follows different power-law functions of separation in the two regimes. He identified the transition separation with the typical Jeans length in the molecular cloud.  However, we show that this is valid only for special cases.  In general, the transition separation depends on the volume density of stars, the depth of the star-forming region, the volume-filling nature of the stellar distribution, and on the parameters of the binaries.  Furthermore, the transition separation evolves with time.  We also note that in young star-forming regions, binaries with separations greater than the transition separation may exist, while in older unbound clusters which have expanded significantly, the transition contains a record of the stellar density when the stars formed.  We then apply these results to the Taurus-Auriga, Ophiuchus, and Orion Trapezium star-forming regions.  We find that while the transition separation in the Taurus-Auriga star-forming region may indicate a typical Jeans length, this is not true of the Orion Trapezium Cluster. We caution against over-interpreting the mean surface density of stellar companions; while Larson showed that Taurus-Auriga is consistent with the stars having a fractal large-scale distribution we show that Taurus-Auriga is also consistent with stars being grouped in non-hierarchical clusters.  We also argue that to make a meaningful study of the stellar distribution in a star-forming region requires a relatively complete stellar survey over a large area. Such a survey does not currently exist for Ophiuchus. Finally, we show that there is no evidence for sub-clustering or fractal structure in the stars of the Orion Trapezium Cluster.  This is consistent with the fact that, if such structure were present when the stars formed, it would have been erased by the current age of the cluster due to the stellar velocity dispersion. ",
        "introduction": "\\label{introduction}  Stars generally do not form in isolation.  Instead, on small scales, they frequently form as members of bound binary or higher-order multiple systems (e.g. Duquennoy \\& Mayor 1991; Mayor et al. 1992; Fischer \\& Marcy 1992; Ghez, Neugebauer, \\& Matthews 1993; Leinert et al. 1993, Simon et al. 1995), while on larger scales they are often members of associations or clusters of stars (e.g. Gomez et al. 1993; Lada, Strom, \\& Myers 1993; Zinnecker, McCaughrean, \\& Wilking 1993). Studying the clustering properties of stars on different length scales may help to determine what processes are involved in their formation.  Gomez et al. \\shortcite{GHKH93} found that the pre-main-sequence stars in the Taurus-Auriga molecular cloud are not randomly distributed, but instead are in small associations of $\\sim 15$ stellar systems within radii of $\\sim 0.5-1.1$ pc.  As one method of analysing the spatial distribution of stars, Gomez et al. determined the two-point angular correlation function and found that it could be represented by a single power-law over separations from $0.005$ to $5$ pc, implying that stars are clustered self-similarly. However, they also found weak evidence that two-point angular correlation function may be better represented by two different power laws with a break at $\\approx 0.05$ pc.  Using data from searches for binary companions to pre-main-sequence stars in the Taurus-Auriga molecular cloud, Larson \\shortcite{Larson95} extended the two-point angular correlation function to smaller separations than Gomez et al. \\shortcite{GHKH93} and demonstrated that, indeed, there is a break at $\\approx 0.04$ pc.  Rather than using the standard two-point angular correlation function, Larson used the closely-related mean surface density of companions (MSDC) (see Section \\ref{MCSD}).  The MSDC has the advantage that no normalisation is required, whereas the two point correlation function must be normalised by the average density in the survey area which can be difficult to determine if the stars are clustered.  Larson \\shortcite{Larson95} found that, for stars in the Taurus-Auriga molecular cloud, the MSDC has a power-law slope of $\\approx -0.6$ on large scales, but steepens below $\\approx 0.04$ pc with a slope of $\\approx -2$ on small scales.  The fact that a break occurs indicates that a single scale-free process is not responsible for the formation of stars on both scales.  The power-law slope of $\\approx -0.6$ on large scales is due to the clustering of stellar systems that Gomez et al. \\shortcite{GHKH93} studied.  Furthermore, Larson pointed out that a power-law slope of $-0.6$ means that the number of stars within an angular distance $\\theta$ of an average star increases as $\\theta^{1.4}$ and, thus, the distribution of stars on this scale can be described as a fractal point distribution with dimension 1.4. Larson identified the power-law slope of $-2$ for small angular separations with the distribution of binary separations, since stellar pairs closer than $0.04$ pc in Taurus-Auriga are typically mutually bound. However, the power-law slope of $\\approx -2$ is not due to a fractal distribution.  Rather, it results from the fact that the frequency distribution of binary separations is roughly uniform in log-separation \\cite{DuqMay91}. Finally, Larson noted that the length scale of $\\approx 0.04$ pc is essentially equal to the typical Jeans length in the Taurus-Auriga molecular cloud.  Thus, Larson associated the location of the break in the MSDC with the Jeans length, speculating that companions with separations smaller than this formed due to the fragmentation of a single collapsing molecular cloud core, while on larger scales stars are grouped self-similarly due to hierarchical structure in the progenitor molecular clouds.  Following Larson's analysis of the Taurus-Auriga star-forming region (SFR), Simon \\shortcite{Simon97} considered the spatial distribution of stars in the Ophiuchus and Orion Trapezium regions. As with Taurus-Auriga, a break was found in the MSDC for each region. On small scales, both Ophiuchus and the Orion Trapezium could be fit by power laws with slopes of $\\approx -2$.  On large scales, flatter power laws were required of $-0.5\\pm0.2$ for Ophiuchus and $-0.2\\pm0.2$ for the Orion Trapezium.  However, the break between the two regimes was found to occur at $\\approx 400$ AU for the Orion Trapezium and $\\approx 5000$ AU for Ophiuchus, compared to $\\approx 10000$ AU (taking the mean of Simon's and Larson's results) for Taurus-Auriga.  Simon concluded that all three SFRs had similar distributions of binary separations and similar fractal structure on large scales, but that the location of the break seemed to depend not only on the Jeans length, but also on the stellar density of the SFR.  Finally, Nakajima et al. \\shortcite{NTHN98} considered the MSDC of stars in the Orion, Ophiuchus, Chamaeleon, Vela, and Lupus star-forming regions.  Again, for those regions where the survey data extends to small enough separations, they find a break in the MSDC with a power-law slope of $\\approx -2$ on small scales and power-law slopes ranging from $-0.15$ to $-0.82$ on large scales.  The location of the break was also found to vary from a minimum of $\\approx 1000$ AU to a maximum of $\\approx 30000$ AU.  Nakajima et al. also considered the nearest-neighbour distributions for each of the regions and found that when the nearest-neighbour distribution could be fit well by a Poisson distribution, the MSDC had a power-law index close to zero on large-scales, while when the nearest-neighbour distribution was broader than the Poisson distribution, the MSDC had a large, negative power-law index. They interpreted this as evidence that the MSDC may indicate a star formation history in the region rather than the presence of self-similar spatial structure; if the stars have a range of ages, the older stars typically will be more dispersed than the younger stars resulting a spread in the distribution of separations of nearest neighbours and a range of stellar surface density which provides the slope of the large-scale MSDC.  Motivated by these papers, we make a careful study of the interpretation of the mean surface density of companions (MSDC) of star-forming regions.  Amongst other goals, we wish to determine the relationship of the break between the binary and large-scale regimes to the Jeans length and the stellar density in star-forming regions.  We also want to determine how sensitive the MSDC is to detecting sub-structure in a stellar distribution and, when detected, what can be said about the form of the sub-structure (e.g. whether the sub-structure is self-similar or not) and how robust the result is.  In Section \\ref{MCSD} we consider the calculation of the MSDC function, handling of survey boundaries, and the results for simple stellar distributions (binaries, global density profiles, clusters, and fractals). In Section \\ref{posbreak} we derive the dependencies of the break between the binary and large-scale regimes, and show that the separation at which the break occurs can only be identified with the Jeans length in special cases.  We also indicate how the MSDC of SFRs is expected to evolve with time. Based on these results, we reconsider the Taurus-Auriga, Ophiuchus, and Orion Trapezium star-forming regions in Section \\ref{application}. Finally, we present our conclusions in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have studied the interpretation of the mean surface density of companions (MSDC) as a function of separation $\\Sigma_{\\rm com}(\\theta)$ in star-forming regions.  We have shown how the power-law slope of $\\approx -2$ for binaries is due their flat distribution of periods in the logarithm of separation, and have considered the MSDC of various global density profiles, sub-clusters and self-similar distributions. We emphasise that simply because a power-law slope can be fit to a particular MSDC, it does not mean that the stellar distribution is self-similar or fractal. We have also demonstrated the effects of survey boundaries on the calculation of $\\Sigma_{\\rm com}(\\theta)$.  Several methods of attempting to avoid boundary effects were considered, all of which provide a full correction in the case that there is no large-scale stellar density gradient across boundaries, but none of which give a perfect correction when there are such large-scale gradients.  Of these, we recommend Method 5, since it allows the maximum range of separations to be studied, does not discard any information, and is simple to use for surveys with irregular boundaries.  Even in the case of a uniform stellar distribution, the improper handing of boundaries results in the $\\Sigma_{\\rm com}(\\theta)$ having a significant slope for separations greater than $\\approx 1/50$ of the survey area's dimensions (i.e.~using Method 1).  Larson \\shortcite{Larson95} associated the separation at which a break in $\\Sigma_{\\rm com}(\\theta)$ between the binary regime and the large-scale regime occurs with the Jeans length in the Taurus-Auriga star-forming region (SFR).  However, we show this transition separation may only be associated with the Jeans length in special cases, and that the transition separation does not necessarily give the maximum binary separation.  In general, the break occurs at the separation where the mean surface density of {\\em binary} companions is equal to the mean surface density of {\\em non-binary} companions (the latter of which may be physically close, or simply chance projections). Thus, typically, the break occurs at smaller separations for SFRs with higher stellar surface densities (as observed by Simon \\shortcite{Simon97} and Nakajima et al. \\shortcite{NTHN98}). In turn, the surface density of non-binary companions depends on the parameters of the binaries, the volume density of stars, the volume-filling factor of the stellar distribution and, in general, the depth of the star-forming region.  The transition separation between the binary and the large-scale regimes also evolves with time.  Due to a stellar velocity dispersion, initial structure is erased and the surface density of stars in an unbound region generally decreases.  This effect begins at the smallest scales, extends to larger scales with time, and results in the transition separation increasing with time.  Finally, the transition between the binary and the large-scale regimes may allow a truncation of binaries at large separations to be detected, especially in old clusters that were much denser when the stars were formed and have since expanded.  In such cases, this provides a record of the stellar density when the stars first formed.  In summary, the transition separation may be associated with the Jeans length only if the star-forming region is young enough that initial structure has not been erased, and if the SFR is `optically thin' in the sense that projection effects due to the depth of the SFR do not affect the transition separation.  The latter is true if the volume-filling factor of the SFR is low (e.g. the SFR is composed of widely separated clusters consisting of only a few stars ($\\sim 10$), or if the stars have a fractal distribution with dimension $\\simless 1.5$).  This is the case for the Taurus-Auriga SFR, which explains the good agreement between between the transition separation and the Jeans length found by Larson \\shortcite{Larson95}, but it is not the case for the Orion Trapezium Cluster.  It is important when studying the large-scale spatial distributions of star-forming regions to obtain the most complete sample of stars over the largest area possible.  The lack of such data for the Ophiuchus SFR makes an attempt to study its large-scale spatial distribution of little use at this time.  For the Taurus-Auriga and Orion Trapezium SFRs, the current data makes a meaningful study of their large-scale stellar distribution possible. For the Taurus-Auriga SFR, Larson \\shortcite{Larson95} fit the large-scale MSDC with a power-law slope that implied a fractal stellar distribution. However, this is not the only possible interpretation; the data can be equally well fit by assuming the stars are formed primarily in randomly-distributed clusters of stars.  For the Orion Trapezium SFR, we find that the MSDC is consistent with the stars simply being distributed according to a surface density that decreases with radius; there is no evidence for sub-structure (either fractal or sub-clusters) in the stellar distribution.  We also demonstrate how upper limits can be placed on how much sub-clustering is present, and note the the sensitivity of the MSDC to detecting sub-structure appears to be slightly less than that of the human eye. The results for the Orion Trapezium SFR are consistent with the fact that if structure were present when the stars formed, it would have been erased by the current time due to the stellar velocity dispersion.  Binaries in the Taurus-Auriga and Orion Trapezium SFR are roughly consistent with an MSDC with a power-law slope of $\\approx -2$.  However, we point out that comparing power-law indices derived from the slope of the MSDC in the binary regime is not the best way to compare the distribution of binary separations between stellar populations since any structure or deviation from a true power-law may easily be missed.  In the centre of the Orion Trapezium SFR, we find very weak evidence that there may be a deficit of binaries with separations $\\simgreat 500$\\,AU\\@.  Such a deficit may be caused by the disruption of wide binaries by single-binary star encounters.  Finally, in view of our studies of the Taurus-Auriga and Orion Trapezium SFRs, we emphasise caution when interpreting the MSDC.  Rather than attempting to characterise star-forming regions simply by fitting power-laws to $\\Sigma_{\\rm com}(\\theta)$, it is more instructive to also consider the global stellar distribution (e.g. a radial surface density profile) and to compare the MSDC to those of model stellar distributions to determine the robustness of any conclusions.  Alternatively, rather than just considering the MSDC (or, equivalently, the two-point correlation function), correlation functions of higher order (three and four-point correlation functions) and/or the nearest-neighbour distribution can be used to differentiate between non-hierarchical and hierarchical structure. The use of higher-order correlation functions is common in studying the large-scale structure of the universe \\cite{Peebles80}.  The nearest-neighbour distribution has been used by Nakajima et al. \\shortcite{NTHN98} to argue that the power-law slope of an MSDC on large scales may indicate a stellar age spread rather than the presence of hierarchical structure.  However, while an age spread does help explain their results, we argue that their results do not exclude the possibility that stars form in hierarchical structures. More work is required on this topic."
    },
    "9804/astro-ph9804012_arXiv.txt": {
        "abstract": "The diffuse X-ray emission from the thin disk surrounding the Galactic mid-plane (the so-called Galactic ridge) was measured with {\\it RXTE} PCA in order to determine the spatial extent, spectral nature, and origin of the emission. Spatial examination of the diffuse emission in the central $30^\\circ$ of the plane in Galactic longitude reveals the presence of two components: a thin disk of full width $ \\lesssim 0^\\circ \\!.5$ centered roughly on the Galactic mid-plane, and a broad component which can be approximated as a Gaussian distribution with FWHM of about $ 4^\\circ$. Assuming an average distance of 16~kpc to the edge of the galaxy, a scale height of about $ 70$~pc and 500~pc is derived for the thin and broad disk components, respectively. Spectral examination of the emission clearly reveals the presence of a hard power law tail above 10~keV and an emission line from He-like iron, indicating both thermal and possibly non-thermal origins for the diffuse emission. The averaged spectrum from the ridge in the $3-35$~keV band can be modelled with a Raymond-Smith plasma component of temperature $\\sim 2-3$~keV and a power law component of photon index $\\sim 1.8$. Based on this finding, we argue that the temperature of the hot phase of the Interstellar Medium (ISM) is less than the previously reported values of $5-15$~keV.  Motivated by the similarities between the characteristics of the thermal component of the Galactic ridge emission in our model and the thermal emission from supernova remnants (SNRs), we discuss the origin of the thermal emission in terms of a population of SNRs residing in the Galactic disk. We find that a SN explosion rate of less than 5 per century is adequate to power the thermal emission from the ridge. The origin of the emission in the hard X-ray band modelled by a power law remains uncertain. Possible contributions from non-thermal bremsstrahlung of cosmic ray electrons and protons, inverse Compton scattering  of energetic electrons from ambient microwave, infrared, and optical photons, non-thermal emission from SNRs, and emission from discrete X-ray sources are discussed. We speculate that bremsstrahlung of accelerated electrons and protons in SNR sites can play a significant role in producing the hard tail of the spectrum. Moreover, their collisional losses can play a major role in the ionization of the ISM. ",
        "introduction": "A galaxy can be well described by an ecosystem. There is an intimate relationship, much like a symbiosis, between the discrete components of the galaxy such as stars, and its interstellar medium (ISM). The ISM provides the foundation for the birth of a new generation of stars, while at the same time it is enriched by the remains of the older generations and their byproducts during their life cycle. Hence, it is natural to expect that the understanding of formation and evolution of galaxies is closely related to the understanding of their ISM.  The ISM of the Milky Way has been found to have many components. Magnetic fields and cosmic ray gas compose the relativistic fluid, while the gaseous phase consists of both ionized and neutral components. The hot ionized medium observable in UV and X-ray has a temperature above $\\sim 10^5$~K and is composed of hot coronal gas heated by supernova shocks. A good portion of the energy of the ISM resides in this component. The warm ionized medium (e.g. HII, planetary nebulae) is visible in ${\\rm H\\alpha}$, UV, and optical, and has a temperature as high as $10^4$~K. The neutral atomic gas appears to have both cold ($<100$~K; e.g. HI~clouds) and warm (100~K~$<T<$~8000~K) components. Molecular clouds compose the self-gravitating gaseous phase. Finally the ISM (with the exception of its hot phase) is filled with dust, particles typically a few tenth of microns in size and visible through their infrared radiation. Since the peak emission of each component arises in different wavebands, multi-wavelength observations from radio to $\\gamma$-rays are needed to probe the ISM.  Diffuse X-ray emission from our galaxy is a powerful diagnostic of the hot phase of the ISM. It is our purpose here to use this tool to probe the processes that contribute to its energetics and dynamics. The first detection of the X-ray emission from the Galactic disk was achieved by the pioneering rocket experiment of Bleach et al. (1972). They detected excess emission associated with a narrow disk component of angular size $2^\\circ-7^\\circ$. Since then, X-ray emission from the Galactic plane and in particular the ridge (the narrow region centered on the Galactic mid-plane extending approximately to $\\pm 60^0$ in longitude and $\\pm 10^0$ in latitude) has been measured in the past with several satellites (e.g. {\\sl HEAO-1} [2-50~keV]: Worrall et. al. 1982; {\\sl EXOSAT} [2-6~keV]: Warwick et al. 1985; {\\sl Tenma} [2-11~keV]: Koyama et al. 1986; {\\sl Ginga} [2-16~keV]: Yamasaki et al. 1997; {\\sl ASCA} [0.5-10~keV]: Kaneda et al. 1997). The presence of the $6.7$~keV iron line in the spectrum discovered with {\\sl Tenma} has motivated the idea that part of the emission below 10~keV is due to a hot optically thin plasma of temperature $5-15$~keV. Because of its high spectral resolution, {\\sl ASCA} has provided the most accurate measurement of the spectrum of the emission to date. The presence of Mg, Si, and Fe K-lines in the spectrum obtained by {\\sl ASCA} suggests that at least part of the emission is of thermal origin. Close examination of {\\sl ASCA} data has also revealed that unresolved, discrete sources are not responsible for the bulk of the emission (Yamauchi et al. 1996; Kaneda 1997). The most recent investigation of the diffuse emission from the Scutum arm region with {\\sl ASCA} (Kaneda et al. 1997) has concluded that the emission below 10~keV has both soft ($kT\\sim 0.8$~keV) and hard ($kT \\sim 7$~keV) thermal components. If indeed the super hot gas ($\\sim 7$~keV) exists in an extended form in the ISM, it is not clear how to explain its confinement to the Galactic disk since its temperature exceeds the gravitational potential of the disk by at least an order of magnitude (Townes 1989). Unfortunately, the two-temperature model does not produce a good fit to the data above 10~keV indicating the presence of additional component(s) at higher energies. Indeed, a hard power law tail has been detected in the hard X-ray/soft $\\gamma$-ray band from observations of the ridge with {\\sl Ginga} and the balloon experiment {\\it Welcome-1} (Yamasaki et al. 1997), and {\\sl OSSE} (Skibo et al. 1997).  In this paper, we present the results from {\\sl RXTE} measurement of the diffuse X-ray emission from the Galactic ridge in the 3-35~keV band. Observations in the hard X-ray/soft $\\gamma$-ray band have usually been complicated by the presence of numerous variable discrete sources, and the fact that the instruments generally have large fields of view and no imaging capabilities, or have imaging capability but no diffuse emission sensitivity. The combination of these factors makes the separation of emission between diffuse and compact sources a difficult task. The advantage of {\\sl RXTE} over previous missions is its small field of view ($1^\\circ$ FWHM) combined with its wide energy bandpass (2-60~keV for the PCA), allowing for the subtraction of the contribution of discrete sources from the diffuse emission spectrum in the hard X-ray band. In addition to reporting on the detection of a hard power law tail in the {\\sl RXTE} data, we also offer an alternative interpretation for the origin of the emission below 10~keV (i.e. instead of a super hot plasma of temperature $5-15$~keV). In agreement with previous studies, our results indicate that the emission is most likely of diffuse origin as opposed to the superposition of discrete sources. However, we present a model in which the X-ray emission is the superposition of both thermal (modelled by a Raymond-Smith plasma) and possibly non-thermal (modelled by a power law) components. We discuss the origin of the thermal component in terms of a population of SNRs residing in the disk. The origin of the power law component remains uncertain. By comparing the spectrum of the diffuse emission in hard X-rays ({\\sl RXTE}) and soft $\\gamma$-rays ({\\sl OSSE}), we find indications that the emission in the two bands are related. We discuss its origin in terms of both discrete hard X-ray sources and radiation mechanisms such as non-thermal bremsstrahlung from cosmic ray electrons and protons, and inverse Compton scattering.   The plan of this paper is as follows. In \\S 2, we describe the observations.  In \\S 3, we present the results of our spatial and spectral analysis of the data. \\S 4 is devoted to the discussion of the results, and their implications for the origin of the emission. Finally in \\S 5, we present our conclusions. ",
        "conclusions": "Our results are summarized as follows: (1)~From the {\\sl RXTE} survey of the Galactic plane, we find that the diffuse emission from the Galactic ridge has two spatial components: a thin disk of width $\\lesssim 0^\\circ\\!.5$, and a broad component with a functional form that can be approximated by a Gaussian distribution of about $4^\\circ$~FWHM. Assuming an average distance of 16~kpc to the edge of galaxy, this translates to a height of 70~pc and 500~pc for the thin and broad disk, respectively. (2)~A hard power law tail is clearly detected in the spectrum and dominates above 10~keV, implying that the emission in hard X-ray possibly has non-thermal origin. On the other hand, the detection of the emission line from He-like iron in the spectrum (and also lower energy lines from Mg, Si, S in the spectrum from {\\sl ASCA}) motivates the idea that part of the emission below 10~keV has thermal origin. (3)~The spectrum in the $3-35$~keV band can be well modelled with a Raymond-Smith plasma component of $2-3$~keV plus a power law component of photon index $\\sim 1.8$. The spectrum above 10~keV simultaneously fitted with {\\sl OSSE}'s spectrum can be fitted with a power law of photon index $2.3\\pm0.2$. The change in the power law slope at lower energies hints at the possibility that the power law either flattens or gradually attenuates below some energy between 10-100~keV. (5)~The characteristics of the thermal component of the diffuse emission resembles that of the SNRs. From this interpretation, we calculate that a SN explosion rate of less than 5 per century is adequate to power the thermal emission from the ridge. (6)~The origin of the emission in the hard X-ray band modelled by a power law is uncertain. While unresolved discrete sources are expected to contribute, the bulk of the emission is expected to be of diffuse origin. One possibility is the non-thermal bremsstrahlung of electrons and protons which may have been accelerated in the SNR sites. Based on the empirical modeling of the data in a wide energy band, we speculate that the power law either flattens or gradually attenuates at lower energies. This lessens the large power injection required to explain the hard power law tail via either electron or proton bremsstrahlung. It will also make it consistent with the power injected to the Galaxy via supernova explosion. (7)~The hydrogen ionization rate implied from our model is $3.3 \\times 10^{-15}n_H^{-1}\\,{\\rm (atom \\, s)}^{-1}$ averaged over the central ridge. Generally, values of $10^{-15}-10^{-14}\\, {\\rm (atom \\, s)}^{-1}$ are expected in the ISM. This is an indication that the collisional losses associated with the bremsstrahlung radiation of energetic particles that produce the hard power law tail are sufficient to explain the ionization rates observed in the ISM.  Finally, we expect that simultaneous {\\sl RXTE}/{\\sl OSSE} observations of the diffuse emission from the Galactic ridge will allow the exclusion of hard discrete sources from the spectrum and will tightly constrain the hard tail slope that dominates in the hard X-ray/soft $\\gamma$-ray band. This will enhance our understanding of the origin of the power law tail and the energetics of ISM."
    },
    "9804/astro-ph9804068_arXiv.txt": {
        "abstract": "\\footnotemark{} \\footnotetext{Table 1 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr(130. 79.128.5) or via http: //cdsweb.u-strasbg.fr/Abstract.html.}    Spectroscopic observations of a sample of 73 very luminous IRAS galaxies ($\\rm log(L_{IR}/L_{\\odot})\\ge11.5$ for $\\rm H_{0}=50km s^{-1} Mpc^{-1}$,$\\rm q_{0}=0.5$) from the 2Jy redshift survey catalogue were carried out using the 2.16m telescope at the Beijing Astronomical Observatory.  The observational data, including the optical images (extracted from Digital Sky Survey) and spectra for these galaxies, are presented in paper I (Wu et al., 1998). In this paper, we give the spectral and morphological classifications for these very luminous IRAS galaxies (VLIRGs). We show that about 60\\% of VLIRGs exhibit AGN-like spectra (Seyfert 1s, Seyfert 2s, LINER-like galaxies). This fraction goes up to 82\\% for the ultraluminous IRAS galaxies (ULIRGs) subsample ($\\rm Log(L_{IR}/L_{\\odot}) \\geq 12.0$). 56\\% of the VLIRGs show strong interaction or merging signatures; this fraction rises to 91\\% for the ULIRGs. These statistical results strongly suggest that interaction triggers nuclear activities and enhances the infrared luminosity. We find that LINER and a mixture type which have optical properties of both HII galaxies and LINERs could be at the transition stage from infrared luminous HII galaxies to AGNs; their main energy production is from starbursts as well as AGNs.  Both infrared luminosities and $\\rm H\\alpha$ equivalent widths increase dramatically as nuclear separations between VLIRGs and their nearest neighbors decrease. There is little doubt that strong starbursts happen in the nuclei of VLIRGs. Assuming class 0 as advanced merger, we construct a simple merger sequence, from morphological classes 1 to 4 (with near or far companions), to class 5 and 6 (interacting pairs and mergers) and then to class 0 (isolated galaxies). Along this sequence, VLIRGs evolve from HII galaxies to AGNs. ",
        "introduction": "Very luminous infrared galaxies (VLIRGs) , especially ultraluminous infrared galaxies (ULIRGs) have bolometric luminosities comparable with quasars and dominate the top end of the galaxy luminosity function (Soifer, et al., 1987). Many theoretical and observational studies (Sanders et al., 1988; Norris et al., 1988; Struck-Marcel and Tinsley, 1978; Rieke et al., 1980; Norman and Scoville, 1988; Rieke, 1988; Condon et al., 1991; Leitherer and Heckman, 1995) show that the high infrared luminosities are associated with the phenomena of active galactic nuclei (AGNs) and starbursts, rather than the radiation from an old stellar population (Thronson, et al., 1990) and cloud-cloud collision (Harwit, et al., 1987).  The main energy output mechanism for luminous infrared galaxies is still being debated. Both  AGNs and starbursts have been proposed as the main energy provider. Sanders et al. (1988) found that a high proportion of AGNs seems to be present in ULIRGs, but Rigopoulou, Lawrence and Rowan-Robinson (1996) found that the starburst model fits well the radio, sub-millimeter/IR to soft X-ray observations for Sanders' 10 ULIRGs. Recent 3D infrared spectroscopic imaging observation (Kroker et al., 1996) suggests that both Seyfert 1 nucleus and circumnuclear star formation contribute significantly to the total luminosity. It is possible that AGNs and starbursts co-exist in the central region. Recently, Veilleux, Sanders and Kim (1997) reported the  results of a near-infrared search for hidden broad-line regions in a sample of 25 ULIRGs with no broad-line regions in the optical band. They detected broad recombination lines in five objects and all of them are Seyfert 2 galaxies. This observation provides important clue that there could exist evolution connection from nuclear starbursts to AGNs in the central region of luminous infrared galaxies.  In order to learn more about the central region of luminous infrared galaxies, and to understand the possible evolutionary process from starbursts to AGNs and how they are associated with galaxy' interaction/merging, we carried out spectroscopic observations for a large sample of 73 VLIRGs compiled from the 2Jy redshift survey catalogue (Strauss et al. 1990, 1992). The spectra and the optical images were presented in a companion paper (Wu et al. 1998, hereafter paper I). In this paper, we report the spectral classification and analysis for the sample (section 2). The morphological classification and their environmental properties are discussed in section 3. In section 4, we point out a possible spectral evolutionary process of luminous infrared galaxies using the relationship between the spectral and morphological classifications. A possible nuclear and evolutionary model for these galaxies is then constructed. Throughout this paper, we use $\\rm H_{0}=50km s^{-1} Mpc^{-1}$ and $\\rm q_{0}=0.5$. ",
        "conclusions": "We have conducted a spectroscopic survey for a complete sample of VLIRGs selected from the 2Jy redshift survey catalogue. We have studied their optical properties and the spectral and morphological classifications for these objects. Important clues are found concerning the nature of VLIRGs and the possible evolution connection between starbursts and AGNs.  \\begin{enumerate} \\bibitem{}  Spectral classification shows that over half (60\\%,44/73) of VLIRGs present AGN-like spectra. As the infrared luminosity increases, this fraction increases from 56\\% ($\\rm 11.5 \\leq Log(L_{IR}/L_{\\odot}) \\leq 12.0$) to 82\\% ($\\rm Log(L_{IR}/L_{\\odot}) \\geq 12.0$). If we exclude LINER-like galaxies, the fraction of AGNs is 13\\%, 36\\% for the two subsamples, respectively (see table 2 a-c).  \\bibitem{} From HII galaxies to AGNs, there is no clear boundary in the diagnostic diagrams (see Figs. 1-3). We classified galaxies with both HII and LINER spectral features as a mixture type. Evidences show that this type of VLIRGs has similar properties as LINERs. We combine the LINERs and the mixture type into a single class as LINER-like galaxies.  \\bibitem{}  Statistics show that there exist correlations among three dust indicators: E(B-V), EW(NaID) and C6563/C4861 in the nuclei of HII galaxies, but not in the nuclei of LINER-like galaxies and Seyfert 2s. This indicates complex nuclear structures of AGNs. The line and continuum emission could come from different nuclear or circumnuclear regions. Seyfert 1's show less reddening.  \\bibitem{}  As separations between VLIRGs and their nearest neighbors decrease, infrared luminosities and  $\\rm H\\alpha$ equivalent widths increase significantly. There exists a character separation, about 10kpc, in ULIRGs, below which the infrared luminosity becomes $\\rm \\ge 2\\times10^{12} L_{\\odot}$. The histogram of the relative velocity shows that many of our sample galaxies have even smaller separations, which provides the favorable condition for triggering starbursts, and hence the high infrared luminosities for theses objects.  \\bibitem{} It is clear from the relationship between spectral and morphological types that the interaction between galaxies triggers massive starbursts as well as nuclear activities. VLIRGs have circumnuclear starburst regions around central AGN. As the spectra evolve from HII-like to AGN-like, the dominant energy output mechanism for VLIRGs changes from starbursts to AGNs.  \\end{enumerate}"
    },
    "9804/astro-ph9804297_arXiv.txt": {
        "abstract": "We present new determinations of bolometric corrections and effective temperature scales as a function of infrared  and optical colors, using a large database of photometric observations of about 6500 Population II giants in Galactic Globular Clusters (GGCs), covering a wide range in metallicity (--2.0$<$[Fe/H]$<$0.0). \\par\\noindent New relations for BC$_K$ {\\it vs} (V--K), (J--K) and BC$_V$ {\\it vs} (B--V), (V--I), (V--J), and new calibrations for T$_{eff}$, using both an empirical relation and model atmospheres, are provided. \\par\\noindent Moreover, an empirical relation to derive the R parameter of the Infrared Flux Method as a function of the stellar temperature is also presented. ",
        "introduction": "A global test of stellar evolutionary models requires a direct comparison between theoretical tracks and observations for stars spanning a wide range in stellar parameters, such as temperature, luminosity and metallicity. In order to achieve these goals at least two fundamental ingredients are needed: \\par {\\it i)~} a complete and homogeneous database of photometric observations; \\par {\\it ii)~} a suitable set of transformations between observables and absolute quantities.  GGCs are the best empirical laboratory to obtain complete and homogeneous spectrophotometric information on Pop. II stars over a wide range of metallicities.  Rewieving the published works on the transformations to the absolute plane (see Sect.3), that is bolometric corrections (BCs) and temperature scales as a function of different colors, it is easy to see that very often these calibrations are not based on a complete and homogeneous set of data, spanning a wide range of stellar parameters. Moreover, {\\it ad hoc} correction factors are usually adopted to take into account for example possible systematic differences between different photometric systems and/or different assumptions for the reference solar quantities, for the adopted model atmospheres, and for different laws to extrapolate the data towards the UV/IR ranges.  Such a scenario indicates that any calibration in the absolute plane can hardly be fully self--consistent as it always depends on the adopted transformations and, more crucial,  the residuals  among different scales are very rarely  linear with the involved parameters.  In order to improve the available determinations of the bolometric corrections and temperature scales, we use here our IR photometric database on GGC stars combined with available optical data from the literature to calibrate new, independent, (hopefully) self--consistent transformations, particularly useful to study the red stellar sequences in Pop II stars.  In Sect.2 we present the complete database used in our analysis which includes about 6500 RGB and HB stars in a sample of 10 GGCs observed in both optical and near IR bands. In Sect.3 we derive the transformations from observed magnitudes and colors to absolute quantities, such as bolometric corrections and effective temperatures. All the results are listed in Table 3. In Sect.4 we compare the inferred scales with existent ones and we give a fully empirical calibration of the R parameter of the Infrared Flux Method (IRFM) as a function of the effective temperature. Schematic conclusions are eventually presented in Sect.5. ",
        "conclusions": "The main results obtained from our analysis can be summarized as follow: \\begin{itemize} \\item By exploiting the use of  a large photometric database of Pop II stars in GGCs we derived  new relations for: BC$_K$ {\\it vs} (V--K) and (J--K), BC$_V$ {\\it vs} (B--V), (V--I) and (V--J), to infer empirical bolometric corrections from observed colors. \\item By making use of both an empirical relation and model atmospheres we calibrated two different scales to infer reliable stellar effective temperatures in the range 3000 -- 7500 K. \\item We also calibrated an empirical relation between the R parameter of the IRFM and the stellar temperature. \\end{itemize} All these relations are summarized in Table 3, where for a given BC$_K$ the corresponding colors and BC$_V$ in two different metallicity regimes and T$_{eff}$ can be read. These relations should be the most suitable to calibrate the red stellar sequences in the color--magnitude diagrams, like the RGB and AGB in old GGCs."
    },
    "9804/astro-ph9804318_arXiv.txt": {
        "abstract": "We present a comprehensive investigation of the radio polarization properties within the theory of natural electromagnetic wave modes in pulsar magnetospheres. Taking into account the curvature of the field lines, aberration effects and magnetic sweep-back we use the relativistic dielectric tensor in the low-density cold plasma approximation and derive the following polarization characteristics, which are in full agreement with the observational findings. Specifically, we demonstrate that {\\bf 1.} The degree of linear polarization decreases with increasing frequency. {\\bf 2.} The degree of circular polarization increases with increasing frequency. {\\bf 3.} At high frequencies ($\\geq$ a few GHz) the degree of linear polarization is correlated to the spin down luminosity $\\dot E$. {\\bf 4.} At high frequencies long-period pulsars exhibit weaker linear polarization than their short-period counterparts. {\\bf 5.} The difference between the refractive indices of the two natural wave modes decreases with increasing frequency which possibly results into a depolarization via superposition. ",
        "introduction": "\\label{intro}  The polarization properties of pulsar radio emission are highly diverse. Nearly all possible polarization states have been detected, from pulsar to pulsar, from frequency to frequency, from pulse to pulse and even within one pulse profile. Nevertheless, some general statements can be made from an observational point-of-view, which must be explained by any plausible theoretical model:  \\begin{itemize} \\item Pulsar radio emission is usually highly linearly polarized at low frequencies (in the following ``low frequencies'' defines frequencies up to about one GHz, ``high frequencies'' corresponds to everything well above one GHz). Towards higher frequencies the radiation tends to depolarize (e.g. \\cite{M71}, \\cite{X96}). \\item The observed polarization position angle (hereafter PPA) is a very stable feature which usually does not depend strongly on frequency. Apart from the occurrence of sudden $\\sim90^\\circ$ jumps, the so-called orthogonal polarization modes (hereafter OPM), the PPA usually follows a characteristic S-shaped curve. This curve can well be modelled by the rotating vector model (hereafter RVM) as proposed by Radhakrishnan \\& Cooke (1969). It is therefore thought to reflect the geometry of the radiating field lines. \\item The sum of many (typically at least a few hundred) single pulses usually converges towards a characteristic average profile. In contrast, the individual pulses show a high variability. One well known single pulse feature is the {\\it quasi-periodic microsecond-structure} (e.g. \\cite{CH77}, \\cite{L98}) observed over a broad frequency interval with no obvious frequency dependence of the quasi-periodicity. It is usually highly linearly and circularly polarized with a roughly constant PPA and often OPMs at the edges. This suggests that, during one micro-pulse, only one of the orthogonal polarization modes is observed. Another feature are the {\\it subpulses}, sometimes drifting in pulse phase from one pulse to the next one. The PPA usually shows a swing during one subpulse, with the swing being stable relative to the subpulse. So when the subpulse drifts, the swing drifts as well, causing a depolarization in the integrated profile (e.g. \\cite{MTH75}). \\item In average profiles and in individual pulses OPMs can be observed (e.g. Stinebring et al. 1984, Gangadhara 1997). The magnitude of these jumps is usually $90^\\circ $ but can also be less sometimes. Sudden orthogonal jumps in the PPA can be found in nearly all pulse phases and even at the same pulse phase in successive single pulses. The OPMs often coincide with changes in the handedness of the circular polarization. Their existence therefore suggests that the radiation preferentially takes two orthogonal elliptical states. \\item Some pulsars show a high degree of circular polarization, preferentially near the centroid of the profile. Sometimes it also sometimes shows a sense reversal near the centroid (\\cite{RR90}). Recent observations show that a few pulsars have a strong increase in their degree of circular polarization towards high frequencies together with a decrease in linear polarization, thus suggesting that a process might be active which transforms linear into circular polarization (\\cite{HKK98}). \\item At low frequencies there seems to be no significant relation between the average degree of polarization and any other pulsar parameter. At higher frequencies, however, correlations have been found between the degree of polarization and the spin down luminosity $\\dot E$ respectively the surface acceleration potential (\\cite{M81}, \\cite{X95} and \\cite{HKK98}). Pulsars with a high $\\dot E$ (respectively a short period and a large period derivative) show a higher average polarization at high frequencies than those with a lower $\\dot E$. This suggests that for these pulsars a possible depolarization process becomes important at higher frequencies than for the other pulsars. Closely connected to this is an anti-correlation between the degree of polarization at high frequencies and the period. Only rapidly rotating pulsars seem to have highly polarized radio emission (\\cite{HKK98}). \\item Several authors estimated the emission altitudes where the observed radiation is produced and studied if there is a frequency dependence of this altitude (Cordes 1978, Rankin 1990, Thorsett 1991, Blaskiewicz et al. 1991, Gil 1991, Phillips 1992, Gil \\& Kijak 1997, von Hoensbroech \\& Xilouris 1997, Kramer et al. 1997). Various independent methods were used and all authors came to the conclusion that the radio emission originates from a region close to the pulsar surface at a few percent of the light cylinder radius (at least for ``normal'' -- non millisecond --  pulsars). The frequency dependence of the emission height is thought to be small, if it exists at all.   \\end{itemize}  All these observational facts have to be explained within a general model for the polarization of pulsar radio emission. Basically, there are two aspects which must be considered. One is the polarization characteristic of the emission process itself, and the other one is the influence of the magnetosphere on the propagating radiation. In this paper we concentrate on the possible role of propagation for the polarization because this can be considered independently of the possible radiation mechanisms.  The magnetosphere of a pulsar is filled with a plasma which is streaming relativistically outwards along the open field lines. A number of articles has already been published on propagational effects in this plasma.  Cocke \\& Pacholczyk (1976) considered the effect of Faraday pulsation for quasi-transverse propagation. Faraday pulsation is the general case of Faraday rotation. It occurs when a polarized beam decays into two elliptical orthogonal modes which propagate with different phase velocities. The PPA of the resulting beam gets rotated and the polarization changes between linear and circular.  \\noindent In a series of papers Melrose \\& Stoneham (1977), Melrose (1979) and Allen \\& Melrose (1982) derived an approximation for the dispersion tensor and calculated the properties of the natural wave modes in the plasma (see also Lyutikov 1998). They assumed that the polarization of the propagating electromagnetic wave follows the shape of the natural wave modes up to a certain radius of limiting polarization (hereafter $R_{\\rm LP}$). The radiation then escapes with the polarization properties of the natural wave modes at $R_{\\rm LP}$. Barnard (1986) places $R_{\\rm LP}$ at the place of cyclotron resonance. For normal pulsars this resonance occurs at a few 10\\% of the light cylinder radius, but for short period pulsars ($P\\leq 0.06 s$) this resonance is outside $R_{\\rm LC}$.  \\noindent Cheng \\& Ruderman (1979)  considered adiabatic walking to be responsible for the variable polarization in subpulses. The polarization properties were derived for different emission mechanisms and plasma compositions.  \\noindent Harding \\& Tademaru (1979, 1980, 1981) presented numerical calculations on the propagation of linearly polarized pulses through a shearing plasma. They show that this can account for micro-structure, circular polarization and rotation of the PPA.  \\noindent Barnard \\& Arons (1986) and Arons \\& Barnard (1986) have calculated the dispersion relations for the X- and the O-mode in an ultra-relativistic, one-dimensional plasma for distribution functions in $e^-$ and $e^+$ (the O-mode has its electric field vector in the plane of curvature of the magnetic field, the X-mode perpendicular to it). The X-mode is purely linearly polarized and propagates easily through the magnetosphere whereas the O-mode can have some circular contributions and follows the bending of the field lines. This causes a separation of the OPMs (\\cite{McK97}).  \\noindent Calculations concerning the propagation characteristics of the various natural modes have also been presented by Beskin et al. (1993). Using the limit of infinite magnetic field strength they conclude that only the electromagnetic X-mode can propagate freely through the magnetosphere. The angular dependence of the properties of the modes and the influence of the cyclotron resonance are derived. This resonance becomes only significant, if the secondary particle production rate is high ($\\ge 10^4$ per primary particle) thus leading to a much higher plasma density. The different mode properties also account for the circular polarization often observed in core components.  In this paper we extended the ideas proposed by Melrose \\& Stoneham(1977), Melrose (1979) and Allen \\& Melrose (1982). We derive the properties of the natural wave modes throughout the magnetosphere with special respect to the angle between the propagating wave and the magnetic field. We make predictions for the qualitative dependence of the polarimetric properties with frequency and various pulsar parameters. These predictions are then compared with the observations. ",
        "conclusions": "We have used the approximation of the dielectric tensor for the low frequency limit to derive the general properties of the natural wave modes in the radio pulsar magnetosphere. The approximation of the dielectric tensor in the given limit implies that the plasma is sufficiently cold. Following Melrose (1979) this should be true for pulsar magnetospheres.  Electromagnetic waves, which propagate through a plasma at an oblique angle to the magnetic field, adopt two orthogonal natural wave modes with individual indices of refraction. We identify these modes with the familiar orthogonal polarization modes. It was our aim to derive the polarimetric properties of these modes and calculate their dependence on different parameters under special consideration of the various angles involved.  We regard only one natural mode assuming that we see only one mode at a time. This assumption is of course not always true, but the existence of highly polarized emission in single pulses indicates that at least sometimes it is true. If both modes are recorded simultaneously, it will lead to a depolarization of the observed radiation. Our results can therefore be regarded as upper limits for the observed degree of polarization. In the case of a highly linearly polarized natural mode we cannot predict the degree of linear polarization one will observe. But if the mode has only a low level of linear polarization, a low level for the observed radiation is predicted.  The main uncertainty of our calculations is the radius of limiting polarization ($R_{\\rm LP}$). However, as we are mainly interested in the qualitative behaviour of the polarization, we have used a fixed value for $R_{\\rm LP}$ without restriction of generality. For a different value of $R_{\\rm LP}$ the numbers may change but the general behaviour remains. This is even true if one moves $R_{\\rm LP}$ to its possible extremes: the light cylinder radius and just above the emission region.  Our results are in excellent agreement with the observations: {\\bf 1.} The degree of linear polarization decreases towards high frequencies (Fig. \\ref{freq1}). The decrease in linear polarization implies an increase in circular polarization (within this concept). This is in agreement with recent observations which show the existence of some pulsars which have this peculiar polarimetric property (\\cite{HKK98}). {\\bf 2.} Given a set of specific pulsar parameters, a variation of $\\dot P$ leads to a variation in the spin down luminosity $\\dot E$. Figure \\ref{edot} shows that the degree of linear polarization of pulsars with a high $\\dot E$ decreases at higher frequencies than for pulsars with a lower $\\dot E$. Thus a correlation between the polarization at high frequencies and $\\dot E$ should exist. Such a correlation has\tbeen indeed observed. {\\bf 3.} Similar to the above point, we have varied the period, keeping all other pulsar parameters fixed (Fig. \\ref{period}). The degree of linear polarization for long period pulsars is expected to decrease at lower frequencies than for short periods pulsars. This should give an inverse correlation between the degree of linear polarization at high frequencies and the period. Again, such an inverse correlation has been observed at a frequency of 5 GHz.  No investigation could be made for the change of the polarimetric characteristics along the different field lines, which are observed during a pulse. The reason is that the qualitative change depends heavily on $R_{\\rm LP}$. For future work it therefore seems to be important to find an estimate for this distance.  The difference between the refractive indices of the two modes decreases with frequency. As this implies a closer propagation of the modes, we expect an increasing superposition. This would lead to a depolarization of the observed radiation at high frequencies. Such a depolarization could account for low abundance of pulsars which show an increasing degree of circular polarization with frequencies.  We note again that our calculations are based on the two assumptions that the background plasma is cold and is dominated by particles with one sign of charge. The obvious advantage of these assumptions -- negligible spread of the energy distribution and no strong pair production -- is their physical and mathematical simplicity. The more it is astounding that we can qualitatively reproduce the complex properties of pulsar radio polarization. In general both assumptions are not widely accepted. Either the energy distribution of the plasma is usually taken to be extended (Arons \\& Barnard 1986; Lyutikov 1998 and references therein) or the secondarily produced electron and positron pairs are expected to dominate the particle content of the pulsar magnetosphere (e.g. Sturrock 1971, Cheng \\& Ruderman 1977 and Daugherty \\& Harding 1982). However the expressiveness of our model deserves further considerations."
    },
    "9804/astro-ph9804182_arXiv.txt": {
        "abstract": "We  numerically investigate chemodynamical evolution of interstellar medium (ISM) in gas-rich disk-disk galaxy mergers in order to explore the origin of fundamental chemical properties of  halo ISM observed in  elliptical galaxies. Main results obtained in this chemodynamical study are the following three.   (1) Elliptical galaxies formed by gas-rich  mergers show steep negative metallicity  gradients of ISM especially in the outer part of galaxies. The essential reason for this is that chemical evolution of ISM in mergers proceeds in such an inhomogeneous way that in the central part of mergers, metal-enrichment of ISM is more efficient owing to radial inflow of metal-enriched ISM during dissipative galaxy merging, whereas in the outer part, metal-enrichment is less efficient owing to  a larger amount of metal-enriched ISM tidally stripped away from mergers. This result provides a clue to the origin of gaseous metallicity gradients of elliptical halo recently revealed by $ASCA$.    (2) Because of inhomogeneous chemical evolution of ISM in mergers, $some$  merger remnants show mean gaseous metallicity discernibly smaller than mean stellar one. The degree of difference in mean stellar and gaseous metallicity in a merger remnant depends  on chemical mixing length, galactic mass, and the effectiveness of supernova feedback.    (3) Elliptical galaxies formed by multiple mergers are more likely to have metal-poor gaseous halo components and steep  gaseous metallicity gradients than those formed by pair mergers. This is principally because a larger amount of less-metal enriched ISM can be tidally stripped away more efficiently from galaxies in multiple mergers.    These three  results demonstrate that dynamical evolution of gas-rich  galaxy mergers, in particular, tidal stripping of less metal-enriched ISM during galaxy merging, greatly determines chemical evolution of ISM of galaxy mergers. These results furthermore imply that recent $ASCA$  observational results concerning mean and radial chemical properties of halo ISM in elliptical galaxies can be understood in terms of chemodynamical evolution of gas-rich galaxy mergers. ",
        "introduction": "Recent observational studies by $ASCA$ ($Advanced$ $Satellite$ $for$ $Cosmology$ $and$ $Astrophysics$) have  revealed a number of fundamental chemical properties of interstellar medium (ISM) of elliptical galaxies, thus have provided valuable information on the formation and evolution of elliptical galaxies (e.g., Awaki et al. 1994; Loewenstein et al. 1994; Matsushita et al. 1994; Mushotzky et al. 1994; Matsumoto et al. 1997; Matsushita et al. 1997). For example, Fe abundance of hot gaseous X-ray halo has been revealed to be appreciably  smaller than that of the stellar component in the host elliptical galaxy (Awaki et al. 1994; Matsushita et al. 1994; Matsumoto et al. 1997; but see Matsushita et al. 1997). This smaller gaseous metallicity is considered to be largely inconsistent with the theoretical prediction of the conventionally  used one-zone chemical evolution model (`` The iron abundance discrepancy problem'') and thus has been extensively discussed in theoretical studies (e.g., Renzini et al.  1993; Fujita, Fukumoto, \\& Okoshi 1996; Arimoto et al. 1997). Furthermore, the hot $X$-ray gaseous halo in elliptical galaxies has been revealed to show strong negative metallicity gradients, which suggests that some physical mechanisms such as cooling flow, gaseous dissipation, galaxy merging, and dilution from external metal-poor gas play a vital role in the formation of the gaseous metallicity gradients (Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). Although these $ASCA$ observational results on mean and radial chemical properties of ISM can provide some diagnostics for any theories of elliptical galaxy formation, only a few theoretical studies have addressed the origin of the above fundamental characteristics of chemical properties of ISM in elliptical galaxies.     The purpose of this paper is to explore the origin of fundamental chemical properties of ISM of  elliptical galaxies recently revealed by $ASCA$. We adopt the assumption that elliptical galaxies are formed by gas-rich disk-disk galaxy mergers and  thereby investigate whether or not the merger model of elliptical galaxy formation can give a plausible explanation for the origin of recent observational results of $ASCA$ on mean and radial chemical properties of ISM in elliptical galaxies. We particularly investigate how the dynamical mixing of chemical components during galaxy merging, which has not been investigated at all in  previous studies, affects mean and radial  chemical properties of ISM of merger remnants. In the present study, the key physical process associated with the origin of metal poor gaseous halo and radial gaseous metallicity gradients in elliptical galaxies is demonstrated to be tidal stripping of less metal-enriched  ISM  during galaxy merging. Based on the present numerical results, we  point out the disadvantages of the commonly used one-zone models in discussing the metallicity of ISM  of elliptical galaxies and stress the importance of dynamical processes of galaxy formation in the chemical enrichment processes of ISM in galaxies. The layout of this paper is as follows. In \\S 2, we summarize  numerical models used in the present study. In \\S 3, we demonstrate how  a number of fundamental $chemical$ properties of ISM in merger remnants are affected by purely $dynamical$ processes of galaxy merging. In \\S 4, we provide a number of implications on chemical properties of ISM in elliptical galaxies. The conclusions of the present study are given in \\S 5. ",
        "conclusions": "There are a growing number of observational evidences which suggest strong radial negative gradients of hot X-ray gaseous halo in elliptical galaxies (Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). Observational study of NGC 4636 (Matsushita et al. 1997) have revealed a factor of $3\\sim4$ difference of  ISM metallicity within $\\sim 7.1 R_{\\rm eff}$. Cooling flows,  stellar metallicity gradients actually existing in ellipticals, the long-term chemical evolution of ISM driven by stellar mass loss or SNIa, and dilution of metal-enriched ISM by external metal-poor gas are considered to be likely explanations for the origin of gaseous metallicity gradients (e.g., Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). These likely explanations are closely associated either with external metal-poor gas or with the later chemical evolution of ellipticals. The present study provides an alternative explanation for  the origin: The ISM  metallicity gradients can be closely associated with intrinsic and dynamical  processes of dissipative galaxy merging at the epoch of elliptical formation. Inhomogeneous and radial-dependent chemical evolution of galaxies  is found to play decisive roles in the formation of negative gaseous metallicity gradients of merger remnants. The present numerical results thus imply that the present-day metallicity gradients of hot ISM halo can contain a fossil record of the past dynamical evolution of ISM of elliptical galaxies at the epoch of their formation.    Negative metallicity gradients derived in the present study are only true for ellipticals with a few Gyr age, primarily because we only solved a few Gyr evolution of chemical properties of ISM but did not solve the later long-term (corresponding to the Hubble time) evolution. Thus, we should investigate the following two points in our future studies in order to confirm whether or not the proposed explanation for the origin of observed metallicity gradients is actually viable for the present-day ellipticals with their ages of  $\\sim$ 10 Gyr in a more quantitative sense. The first is long-term chemical evolution of ISM surrounding merger remnants, which can be  greatly affected by the later and continuous metal-enrichment due to SNIa and stellar mass loss of long-lived stars. As has been demonstrated by several previous studies on long-term hydrodynamical evolution of hot ISM in ellipticals (e.g., David, Forman, \\& Jones 1991), the effectiveness of thermal heating by SNIa (and partly by long-lived stars) determines time-evolution radial flow patters  of hot ISM (e.g., either outflow due to effective  thermal heating of SNIa or inflow due to efficient cooling). Such later gaseous inflow or outflow, which can transfer metal-enriched ISM, can change radial metallicity distribution that is initially formed by gas-rich galaxy merging. Accordingly, chemodynamical effects of the later stellar mass-loss and supernovae on the radial metallicity distribution should be explored more in detail in our future studies. The second is the long-term dynamical evolution of less metal-enriched ISM that is tidally stripped and surrounds merger remnants. The present study predicts that the tidally stripped metal-enriched ISM surrounds the considerably outer part of merger remnants, where external tidal field resulting from neighbor galaxies and  large-scale gravitational potential of cluster or group of galaxies can affect dynamical evolution of the stripped ISM. Later dynamical effects of external tidal force can change drastically initial radial gradients of ISM in merger remnants, thus we should also consider  these in our future studies. Since radial  metallicity gradients of ISM in ellipticals contain valuable information not only on chemical evolution of ISM but also on dynamical evolution of galaxies as a whole, more extensive studies are necessary for the better understanding of the origin of the gradients.    Furthermore,  the difference in mean metallicity between stellar and gaseous components derived in the present merger model can be compared with recent observational results by $ASCA$ (Awaki et al. 1994; Matsushita et al. 1994; Matsumoto et al. 1997) which reveals that the abundance of hot gaseous X-ray halo (Fe) is appreciably  smaller than that of the stellar component in the host elliptical galaxy. These observational results seem to be consistent with results derived  in some merger models of the present study, which imply that metal-poor gaseous components in some elliptical haloes can be formed by gas-rich mergers. The most recent result of Matsushita et al. (1997), however, shows that there is not so large difference between mean stellar metallicity  (0.74 solar) and mean gaseous one ($\\sim$ 1.0 solar) within $\\sim 4.0 R_{\\rm eff}$ of NGC 4636 and furthermore that the gaseous metallicity in the outer part of the halo (for the region $4.7 \\leq R/R_{\\rm eff} \\leq 7.1$) is still smaller ($\\sim 0.37$  solar) than mean stellar metallicity of NGC 4636. Although  metallicity averaged out for the whole region of the gaseous halo in NGC 4636 has not been clarified yet, this  new observational result provides the following implication on inhomogeneous chemical mixing derived in the present study. If the gaseous metallicity averaged out for the whole region of the gaseous halo is really larger than stellar one in NGC 4636, the present numerical result that mean gaseous metallicity can be smaller than mean stellar one in some merger remnants is not consistent  with the observational result. In this case, we should consider either that inhomogeneous chemical mixing is not so efficient in real galaxy mergers as the present study predicts, or that the later  long-term metal-enrichment of ISM resulting from metal-ejection from long-lived stars and supernovae (SNIa) can greatly affect the chemical evolution of ISM after galaxy merging and thus change the difference of  mean stellar and gaseous metallicity in merger remnants. Alternatively, if the gaseous metallicity averaged out for the whole region of the gaseous halo is really  smaller  than stellar one, the present study can provide  a clue to the origin of such  smaller gaseous metallicity of elliptical haloes: Origin of metal-poor gaseous halo in ellipticals can be closely associated with the past dissipative galaxy merging processes. Since  total number of sample X-ray gaseous halo with the mean metallicity estimated precisely is still small, it is safe for us to say, at least now, that future more extensive observational studies and more elaborated theoretical models will assess  the validity of inhomogeneous chemical mixing derived in the present study.  Main results obtained in the present  chemodynamical study are the following three.   (1) Elliptical galaxies formed by gas-rich  mergers show steep negative metallicity  gradients of ISM especially in the outer part of galaxies. The essential reason for this is that chemical evolution of ISM in mergers proceeds in such an inhomogeneous way that in the central part of mergers, metal-enrichment of ISM is more efficient owing to radial inflow of metal-enriched ISM during dissipative galaxy merging, whereas in the outer part, metal-enrichment is less efficient owing to  a larger amount of metal-enriched ISM tidally stripped away from mergers. This result provides a clue to the origin of gaseous metallicity gradients of elliptical halo recently revealed by $ASCA$.    (2) Because of inhomogeneous chemical evolution of ISM in mergers, $some$  merger remnants show mean gaseous metallicity discernibly  smaller than mean stellar one. The degree of difference in mean stellar and gaseous metallicity in a merger remnant depends  on chemical mixing length, galactic mass, and the effectiveness of supernova feedback.    (3) Elliptical galaxies formed by multiple mergers are more likely to have metal-poor gaseous halo components and steep  gaseous metallicity gradients than those formed by pair mergers. This is principally because a larger amount of less-metal enriched ISM can be tidally stripped away more efficiently from galaxies in multiple mergers.     These three  results demonstrate that dynamical evolution of galaxy mergers can greatly affect chemical evolution of ISM of galaxies, which cannot be attained untill both dynamical and chemical evolution of galaxies are solved in an admittedly self-consistent manner. In particular, tidal stripping of less metal-enriched ISM during dissipative galaxy merging is found to play a vital role in determining mean and radial chemical properties of ISM in elliptical galaxies. The present study accordingly implies that the origins of metal-poor gaseous halo and negative metallicity gradients of ISM in  an elliptical galaxy can be  closely associated with gas-rich galaxy merging at the epoch of elliptical galaxy formation."
    },
    "9804/astro-ph9804307_arXiv.txt": {
        "abstract": "We investigate the spectroscopic characteristics of the optical components of Be/X-ray binary systems, using data collected during our seven-year monitoring campaign. We find examples of major changes in the emission line profiles associated with Type II X-ray outbursts, later developing into V/R variability cycles. We show that the time-scales for V/R variability in Be/X-ray transients extend from a few weeks to years and interpret all these changes as due to the presence of global disruptions of the axisymmetric density distribution in the extended envelopes of the Be stars in these systems. The association between X-ray outbursts and V/R variability, the occurrence of very fast changes and the very short quasi-periods of variability displayed by Be/X-ray binaries lead us to conclude that the presence of the neutron star is an important factor affecting the dynamics of the disc-like envelopes. The interaction between the compact companion and the disc would explain the correlation between H$\\alpha$ strength and orbital period recently found. The characteristics of the V/R cycles are, however, mainly independent of the binary parameters. ",
        "introduction": "Be/X-ray binaries constitute the major subclass of massive X-ray binaries, in which  X-ray emission is due to accretion of matter from an early-type  mass-losing star by a compact companion (see Apparao 1994, White et al. 1995, for reviews). Be stars are early-type  non-supergiant stars, which at some time have shown emission in the Balmer lines. Both the emission lines and the characteristic strong infrared excess when compared to normal stars of the same spectral types are attributed to the presence of a cool circumstellar envelope, presumably in the shape of a disc (see Slettebak 1988). The physical reasons which give rise to the disc are unknown, but it is generally believed that the high rotational velocity of Be stars plays an important role, even though it is accepted that some other mechanism(s) must be at work.  Most Be/X-ray binaries have relatively eccentric orbits and the compact companion (in general, a neutron star, but in some cases possibly a white dwarf) spends most of its time far away from the disc surrounding the Be star. Three kinds of X-ray activity are observed (Stella et al. 1986, henceforth SWR):  \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item Persistent low-luminosity ($L_{{\\mathrm x}} \\la 10^{36}$ erg s$^{-1}$) X-ray emission. Some sources (e.g., X Persei) have always been observed in this state.  \\item Periodical (Type I in SWR) X-ray outbursts ($L_{{\\mathrm x}} \\approx 10^{36} - 10^{37}$ erg s$^{-1}$), coinciding with the periastron passage of the neutron star. Type I outbursts have been observed in numerous sources, such as A\\,0535+262 (Motch et al. 1991) and EXO\\,2030+375 (Norton et al. 1994).  \\item Giant (Type II in SWR) X-ray outbursts ($L_{{\\mathrm x}} \\ga 10^{37}$ erg s$^{-1}$), which do not show any orbital modulation. Type II outbursts are normally seen in those sources that also display Type I activity (see Parmar et al. 1989, Finger et al. 1996a for examples). \\end{enumerate}  Be/X-ray binary systems which display outbursts are collectively termed Be/X-ray transients. Most transients (e.g. A\\,0535+26) also show low-luminosity X-ray emission when they are not in outburst, but in systems containing fast-rotating neutron stars, centrifugal inhibition of accretion prevents X-ray emission (SWR) except during outbursts (e.g., 4U\\,0115+634). Type I outbursts occur in series between long periods of X-ray inactivity (or low-luminosity emission), while the onset of Type II outbursts is completely unpredictable. ",
        "conclusions": "We have presented observational evidence showing that global disruptions are frequent in the extended circumstellar envelopes of Be/X-ray binaries. These perturbations are reflected in the asymmetric line profiles normally observed from these systems. V/R ratio variability is observed to occur with typical time-scales ranging from a few days to several years. In at least two cases (the giant outbursts of A\\,0535+26 in February 1994 and 4U\\,0115+634 in December 1995), a major disruption seems to have originated in coincidence with the X-ray outburst. Further evidence of the association between fast changes in the line profiles and X-ray outbursts has been seen in most Be/X-ray transients.  We believe that all these observation suggest that the presence of the neutron star represents a major factor controlling the dynamics of the discs around the Be stars in X-ray binaries. This fact provides an explanation to the correlation between maximum H$\\alpha$ EW and orbital period found by Reig et al. (\\cite{reigb}). The frequent presence of major density perturbations in the envelopes of Be/X-ray binaries introduces a new element of complication in the modelling of these systems. Rather than assuming that the disc is static and homogeneous, new models should take into account the presence of global density waves and explore the possibility that the series of Type I outbursts are caused by the interaction of the neutron star with the regions of enhanced density which these waves generate.   Continued monitoring of Be/X-ray transients and careful optical coverage of future Type II outbursts will provide the only test for these hypothesis."
    },
    "9804/astro-ph9804077_arXiv.txt": {
        "abstract": "We have observed with the IRAM interferometer at two different epochs and simultaneously the two transitions $v=0, J=2 \\rightarrow 1$ and $v=1, J=2 \\rightarrow 1$ of \\dSiO and \\uSiO in Orion IRc2. We have made the first maps of \\dSiO $v=0, J=2 \\rightarrow 1$ emission from Orion. These maps and properties of the \\dSiO spectra attest to maser emission. Our \\uSiO maps show the stable ring of maser spots observed in previous works. Combining our own data with published works we derive that the relative motion between the two ridges of the \\uSiO emission ring is less than about 0.7 AU/yr over a period of 7 years. On the other hand, the weak high velocity maser features observed around 30 \\kms ~seem to move with respect to the stable ring of \\uSiO main emission. Our relative \\dSiO ($v=0$) and \\uSiO ($v=1$) spot maps show that most \\dSiO and \\uSiO emission features are closely related but have not the same spatial extent. We conclude that these masers are not excited in the same gas layers in agreement with pumping models which predict that various $v$ state masers peak in different spatial regions. In addition, our maps of $v=0$ and $v=1$ emission suggest that local line overlaps due to turbulence and high gas temperature do not play a dominant role in the excitation of \\uSiO and \\dSiO, although excitation effects resulting from the overlap of Doppler-shifted ro-vibrational lines may still be significant. ",
        "introduction": "Since the discovery nearly 25 years ago of the SiO molecule in Orion (Snyder \\& Buhl 1974), SiO masers have been observed in the envelopes of hundreds of late-type stars, and in the direction of a few galactic HII regions. Orion remains a unique source of SiO emission because it contains all known SiO isotopic species and because it is the only star-forming region exhibiting a very strong maser in the $v=1, J=2 \\rightarrow 1$ transition. The strong and compact $v=1$ SiO maser is associated with the luminous infrared source IRc2 and is closely connected with the extended and weaker $v=0$ maser emission mapped at 43 GHz (Chandler \\& De Pree 1995) and 86 GHz (Wright et al. 1995). Recent high resolution observations showed that IRc2 is a complex object resolved into four components (Dougados et al. 1993) and that the center of the strong ($v=1$) and weak ($v=0$) SiO maser outflows coincide with the radio continuum source I (Menten \\& Reid 1995, Wright et al. 1995), and is displaced from the center of the molecular hot core (lying to the east of source I). On the other hand, the center of the large-scale high velocity bipolar outflow traced by CO lies roughly $3\"$ to the north of source I.  Considerable efforts have been made to properly model the SiO maser phenomenon in late-type stars. Models include radiative and/or collisional excitation schemes (e.g. Kwan \\& Scoville 1974, Elitzur 1980, Langer \\& Watson 1984, Lockett \\& Elitzur 1992 or Bujarrabal 1994). All models share two general characteristics: $(i)$ they require high volumic densities of order $10^{8}-10^{10}$ cm$^{-3}$; $(ii)$ inversion of the SiO level populations depends on the column density, and maser emission in higher vibrational states peaks at higher values of the column density. On the other hand, there are major observational facts that cannot be explained by any of the present radiative/collisional SiO pumping schemes. In particular, the \"standard\" pumping schemes fail to explain in stars the peculiar distribution of line intensities within a given vibrational state (e.g. Cernicharo \\& Bujarrabal 1993), and fail to explain the absence or weakness of $v=2, J=2 \\rightarrow 1$ emission from Orion and late-type stars (Olofsson et al. 1981 b, Bujarrabal et al. 1996). In fact, line overlaps among transitions of the isotopic species of silicon monoxide are an important addition to radiative/collisional pumping in stars (e.g. Gonz\\'alez-Alfonso \\& Cernicharo 1997), while the line overlap between two near infrared lines of SiO and water explains the weakness of the $v=2, J=2 \\rightarrow 1$ transition (Bujarrabal et al. 1996).  Depending on the relative importance of property $(ii)$ above (namely various $v$ state masers should peak in different spatial regions) with respect to line overlap effects among nearby transitions of SiO and isotopes the spatial distribution of various vibrational transitions should differ or not. Therefore, high spatial resolution and sensitive maps of SiO and isotopes should provide a test of these predictions. With this idea in mind we have compared the spatial distributions of two nearby transitions of \\dSiO and \\uSiO toward Orion IRc2 which contains the strongest SiO source in the sky. In Sect. 2 we present our observations and give details of data reduction. In Sect. 3 we discuss spectral variability and present our maps of \\dSiO and \\uSiO emission from Orion. In Sect.~4 we discuss some properties of the apparent ring of \\dSiO and \\uSiO masers, the relative spatial extents of both species and implications on their excitation. Some conclusions are summarized in Sect. 5. ",
        "conclusions": "\\subsection{Stability of the ring of \\uSiO masers} The regular position$-$velocity pattern of $v=1, J=2 \\rightarrow 1$ emission (Plambeck et al. 1990, Wright et al. 1995) was interpreted by Plambeck et al. as a collection of maser clumps lying in an expanding and rotating disk. This pattern does not change much with time and has been observed with the IRAM interferometer in 1990 (Guilloteau et al. 1992), 1992 (Baudry et al. 1995), and 1995 and 1996 (this work). A similar position$-$velocity pattern was also observed in the $J=1 \\rightarrow 0$ transition of SiO around 43 GHz by Morita \\etal (1992) and Menten \\& Reid (1995). We can estimate the long-term stability of the 86 GHz pattern by measuring the mean separation between the two ridges delineated by the dominant positive and negative velocity features of Orion IRc2. To this end, we have used the 6 different maps made at 86 GHz with the BIMA and IRAM interferometers (Plambeck et al. 1990, Wright et al. 1995, Guilloteau et al. 1992, Baudry et al. 1995, and this work). The general orientation and the mean separation between the two ridges of main SiO features do not seem to evolve with time. The mean separation between these two ridges is $\\approx 0.^{''}165 \\pm 0.^{''}01$; for the uncertainty we have assumed that the 6 independent measurements behaved as gaussian variables. Therefore, any apparent contraction or expansion of the ring, would be less than or of order $0.^{''}01 / 7$yr, namely $\\leq 0.7$ AU/yr at the 480 pc distance of Orion A. On the other hand, stability of the intermediate velocity pattern ($\\approx 0 \\rightarrow 11$ km s$^{-1}$) is not obvious when we compare our 1990 data with the present IRAM maps. The complex shape observed in 1995 or 1996 is not quite similar to that in 1990 (see Fig. 11 in Guilloteau et al. 1992). Such differences cannot be due to relative position errors which are less than about 2 to 5 mas in the \\uSiO maps. These discrepancies seem to agree with the model of Plambeck et al. (1990) which predicts that intensity changes in the intermediate velocity features could cause large position changes of the maser spots in the disk.  \\subsection{Nature of the \\uSiO high velocity features}  The high velocity features lying around $28\\rightarrow 31$ km s$^{-1}$ are located close to the positive velocities of the SiO ring (see e.g. the 29.2 and 29.1 km s$^{-1}$ features in Figs. \\ref{relat28map} and \\ref{relat2829map}). Our observations of 1995 August show that these features are excited in an area similar, although not identical, to that observed in 1995 January by Wright et al. (1995, see their Fig. 1c) for their $30 \\rightarrow 33$ km s$^{-1}$ features. However, 6.5 months later our data show that the $28 \\rightarrow 31$ km s$^{-1}$ features have migrated toward the most positive velocity end of the main SiO emission ridge (see location of the 29.1 km s$^{-1}$ component in Fig.\\ref{relat2829map}). The apparent migration of weak high velocity components is consistent with the model of Plambeck \\etal (1990) where small changes in brightness distribution of extended features may look like rapid motion. Nevertheless, such rapid motions should be confirmed in future maps of \\uSiO emission. These high velocity features could be related to the spectral changes observed for the same components; they could be weakly masing as suggested in Sect. 3.2. It is interesting to note that anomalous gas motion beyond the expansion velocity of the ring of maser clumps could perhaps explain the high velocity components. Such components are reminiscent of the weak features observed in the line wings of \\uSiO emission from late type stars (Cernicharo et al. 1997, Herpin et al. 1998). SiO line wing emission in stars is related to bipolar gas outflows and to pulsations of the underlying star. \\begin{figure} [t] \\begin{center} \\epsfxsize=8.5cm \\epsfbox{figure6vit.epsf} \\end{center} \\caption[]{Comparison of \\uSiO $v=1, J=2 \\rightarrow 1$ and \\dSiO $v=0, J=2 \\rightarrow 1$ spot maps using the \\uSiO feature at 15.6 km s$^{-1}$ as a phase reference in both maps. The epoch of the observations was March 5, 1996 for both isotopes. We have plotted the centroids of the \\uSiO features (open squares and small full circles) together with the main \\dSiO features (dotted circles). The diameter of each circle is proportional to the peak intensity for \\dSiO. The LSR velocity labels are given every 5 and 6 channels for \\uSiO and \\dSiO, respectively.} \\label{relat2829map} \\end{figure} \\begin{figure} [tbh] \\begin{center} \\epsfxsize=8.5cm \\epsfbox{figure6rond.epsf} \\end{center} \\caption[]{Comparison of \\uSiO $v=1, J=2 \\rightarrow 1$ and \\dSiO $v=0, J=2 \\rightarrow 1$ spot maps using the \\uSiO feature at 15.6 km s$^{-1}$ as a phase reference in both maps. The epoch of the observations was March 5, 1996 for both isotopes. We only show the centroids of \\dSiO (dotted circles) and \\uSiO (full circles) emissions. The diameter of each circle is proportional to the peak intensity.} \\label{relat2829mapbis} \\end{figure} \\subsection{Nature of the \\dSiO $v=0$ emission} Time variability as well as notable changes in the \\dSiO line profile (Fig. \\ref{variable29SiO}) are clearly in favour of maser emission. Short-term variability in the excitation of the \\dSiO molecule is suggested by our maps because the mean distance observed in 1996 between the two ridges of emission, $\\approx 0.13\"-0.14\"$, is significantly larger than the $0.11\"$ measured in 1995. This fact indicates non thermal processes in the excitation of \\dSiO. In addition, the relatively high flux density observed in 1995 and 1996 suggests also non thermal emission. The array cannot give the size of the individual features, but we may use the synthesized beamwidth and the observed peak flux density to estimate a minimum brightness temperature. In March 1996 the flux density peaks around 37 Jy and we derive $T_{B} \\geq 2500$ K. This temperature is greatly above the kinetic temperature usually adopted in Orion, $\\approx 60$ K, and is another indication for maser emission although high temperatures would be plausible in a shocked environment. However, for some of the weaker \\dSiO features in the range $0\\rightarrow 10$ \\kms we obtain T$_{B}\\geq 80-100$ K, and \\dSiO could thus be part thermal and part maser.  The actual spatial structure of the \\dSiO $v=0$ emission is complex, and we recall that Fig. \\ref{29SiOrelatmap} does not show all of the emission detectable with the array. By discarding the shorter baselines (Sect. 3.1), we have concentrated our analysis on more compact emission sources. We were not able here to map the \\dSiO $v=0$ emission counterpart to the extended \\uSiO $v=0$ emission seen by Wright et al. (1995).  \\subsection{Comparison of the \\dSiO and \\uSiO spot maps. Implication on excitation mechanisms} The relative distribution of \\uSiO and \\dSiO emission observed in 1996 is shown in Fig. \\ref{relat2829map} and in Fig. \\ref{relat2829mapbis} where, for clarity, we have not given the velocities. The two ridges of \\dSiO $v=0, J=2 \\rightarrow 1$ emission (dotted circles) and \\uSiO $v=1, J=2 \\rightarrow 1$ emission (full circles) are clearly visible on these figures. Our data do not show a complete spatial overlap as well as a close correlation among features of both species although in both cases the positive velocities lie to the NW of the negative features. Another clear difference between \\dSiO and \\uSiO is that there is no obvious pattern for the \\dSiO intermediate velocity features (Fig. \\ref{29SiOrelatmap} or Fig. \\ref{relat2829map}); this could be related to mixed thermal and masing features as suggested in the previous Section.  Fig. \\ref{relat2829map} shows that velocities in the range $ \\approx -6 \\rightarrow -10$ km s$^{-1}$ tend to be found in the same area for both species although it is not possible to make an exact position-velocity pairing of the \\dSiO and \\uSiO features. In Fig. \\ref{relat2829mapbis} the northern ridges of both \\dSiO and \\uSiO seem to be co-aligned and are not coincident. We cannot exclude, however, that uncorrected instrumental effects still affect our relative map and these observations should be repeated using frequent bandpass calibrations as used in 1996. In order to force the spatial coincidence of the main features in both species we have shifted the main \\dSiO features lying around 16 km s$^{-1}$ on top of the \\uSiO features in the range $15 \\rightarrow 17$ km s$^{-1}$. Nevertheless, the \\dSiO negative velocity emission ridge appears well outside the \\uSiO negative velocity ridge. Any rotation of coordinates axis around the 16 km s$^{-1}$ features does not improve the spatial coincidence of both species. Hence, we conclude that both isotopic species are not excited in the same gas layers. This is strengthened by the analysis of our 1995 data which similarly show no spatial coincidence and a smaller distance between the two ridges of emission for \\dSiO than for \\uSiO .  We note that a similar picture also emerges from the 43 GHz interferometric observations made by Morita et al. (1992) in Orion. Although their observations of the \\uSiO $v=2, J=1 \\rightarrow 0$ and $v=1, J=1 \\rightarrow 0$ transitions were not made simultaneously and were less sensitive than here, the mean separation between the two ridges of emission is slightly smaller for $v=2$ ($\\approx 0.12\"$) than for $v=1$ ($\\approx 0.14\"$ in agreement with the separation measured on the 43 GHz map of Menten \\& Reid 1995). The 43 GHz observations, the maps of \\uSiO $v=1$ and $v=0$ emission (Wright et al. 1995) and our 86 GHz maps indicate that different vibrational states of silicon monoxide do show a close connection but do not exactly coincide.  We comment below on possible explanations of the observed similarities without exact co-location of \\dSiO ($v=0$) and \\uSiO ($v=1$) maser sources. First, the silicon monoxide reservoir seems identical for both isotopic species since their large-scale spatial distributions are alike. This is expected if shocks generated in the expanding flow traced by the $v=1, J=2 \\rightarrow 1$ masers enhance the sputtering of silicon which will then react quickly in the gas phase to form both \\dSiO and \\uSiO. Second, differences in the small-scale spatial distributions of \\dSiO and \\uSiO could simply result from differences in the excitation of both species or from different physical conditions within the silicon monoxide cloud. SiO pumping models do not require any isotopic differentiation to obtain \\uSiO, \\dSiO or $^{30}$SiO maser sources. In all cases the general physical conditions are grossly similar for one isotopic species or another apart from the total column densities. On the other hand, all radiative/collisional pumping models show that different $v$ state masers peak in different spatial regions. We believe that this fact, combined with different degrees of saturation in the \\uSiO and \\dSiO masers, is essential to explain the slightly different distribution of \\dSiO ($v=0$) and \\uSiO ($v=1$) maser spots. Our maps of relative \\uSiO and \\dSiO emission also show that some features from both isotopes and with different velocities tend to be excited in the same area. This is observed in the range $-6 \\rightarrow -10$ km s$^{-1}$. Collisional pumping with high temperature ($\\approx 1500$ K) and high molecular hydrogen density ($\\approx 10^{9}-10^{10}$ cm$^{-3}$) provides a range of \\uSiO column densities where both $v=1$ and $v=2$, $J=1 \\rightarrow 0$ masers are excited (Lockett \\& Elitzur 1992). Such a scheme does not apply to our apparently overlapping \\dSiO ($v=0$) and \\uSiO ($v=1$) 86 GHz features because their velocities are not in good agreement. However, further observations should be conducted to investigate the detailed kinematics and stability of the \\dSiO emission.  Line overlap effects among various transitions of silicon monoxide cannot be ignored to explain the excitation of this molecule. First, {\\em local} line overlaps due to turbulence play a role as soon as the velocity dispersion reaches about 5 \\kms. Limiting ourselves to the lower $J$ values, we find that 10 to 15 ro-vibrational transitions of \\uSiO, \\dSiO and \\tSiO overlap within 5 \\kms for $\\Delta v=2$ and 1. If local line overlaps would dominate the excitation of low $J$ rotational levels in Orion, we would expect exact spatial coincidence of the isotopic species. Our \\dSiO and \\uSiO maps contradict this idea. Second, {\\em non-local} line overlap effects as described by Gonz\\'{a}lez-Alfonso \\& Cernicharo (1997) in a non static circumstellar environment are most important. The relative distribution of \\uSiO and \\dSiO emission in our maps is not inconsistent with such non-local line overlaps. In addition, it is also plausible that the overlap between two near infrared lines of water and \\uSiO (Olofsson et al. 1981 b) is an important excitation process of silicon monoxide in Orion.  Analysing the spatial extents of different $v$ state masers is clearly important to better understand the pumping mechanisms of the SiO molecule. This kind of work should be extended to strong stellar SiO-emitters since the physical conditions in late-type stars and Orion are so different. VLBI observations are required in stars in order to make a detailed comparison of the different $v$ emission layers; such observations have been made for the first time to map the $v=2, v=1, J=1 \\rightarrow 0$ lines in W Hya and VY CMa (Miyoshi et al. 1994)."
    },
    "9804/astro-ph9804288_arXiv.txt": {
        "abstract": "We have calculated several representative models of vertical structure of an accretion disk around a supermassive Kerr black hole. The interaction of radiation and matter is treated self-consistently, taking into account departures from LTE for calculating both the disk structure and the radiation field. The structural equations are described in detail, and various approximations are discussed. We have demonstrated that departures from LTE are very important for determining the disk structure, even at the midplane, as well as the emergent radiation, particularly for hot and electron--scattering--dominated disks. We have shown that at least for the disk parameters studied in this paper, NLTE effects tend to reduce the value of the Lyman jump with respect to the LTE predictions, regardless whether LTE predicts an emission or absorption jump. We have studied the effects of various values of viscosity on the model structure and predicted spectral energy distribution. The viscosity is parameterized through a parameter $\\alpha_0$ which describes the vertically-averaged viscous stress, two power-law exponents $\\zeta_0$ and $\\zeta_1$, and the division point $m_{\\rm d}$ between these two forms. The disk structure and emergent radiation is sensitive mainly to the values of $\\alpha_0$, while the other parameters influence the disk structure to a much lesser extent. However, although the detailed shape of the predicted spectrum is sensitive to adopted value of $\\alpha_0$, the overall appearance of the spectrum is quite similar. ",
        "introduction": "Accretion disks around massive black holes have long been the most popular candidates for providing the ultraviolet and soft X-ray flux observed in Active Galactic Nuclei (AGN). Observational evidence is mainly based on the `big blue bumps' seen in the UV (e.g. Shields 1978; Malkan \\& Sargent 1982). However, despite its attractiveness this model faces a number of problems when confronted with multi-wavelength observations. The current situation has been recently summarized by Koratkar (1998) from the observational point of view, and by Blaes (1998) from the theoretical viewpoint. The most pressing problems of the theoretical models are a near absence of observed Lyman discontinuity (first pointed out by Antonucci, Kinney, \\& Ford 1989), the UV/EUV continuum polarization (e.g., Antonucci 1992, Blaes 1998 and references therein), the overall continuum spectral energy distribution (e.g. Laor 1990), and phased optical/UV variability (e.g., Alloin et al 1985).  What is clearly needed is a self-consistent model which would explain all the observed features. Such a goal is still far away, but we feel that the first step towards it is to answer the fundamental question: In view of all current problems, is the accretion disk paradigm still a viable model? In other words, are the observed phenomena truly inconsistent with the accretion disk picture in general, or is the lack of agreement rather a result of inaccuracies or even inconsistencies in the computed models or in the current modeling techniques?  Therefore, we have embarked on a systematic study of these questions. We recognize that there are many possible sources of inconsistencies in the AGN accretion disk modeling. There are essentially two types of problems. First, there are basic physical uncertainties. Among them, the most important is our ignorance of the basic physics of viscous energy dissipation in the AGN disks, so that we are left with a necessity to employ certain {\\it ad hoc} parameters. Although this is certainly a viable approach when nothing else is currently available, we should bear in mind that corresponding models will lack predictive power---we may explain what we see, but we cannot predict it. In any case, when adopting a model based on some chosen set of parameters, one at least has to study carefully the influence of these parameters on the computed model. Other fundamental physical problems are uncertainties of the effects of other structures forming the AGN complex onto the accretion disk; for instance an irradiation of the disk from external sources. Again, in the absence of any detailed theory, we are usually left with a necessity to parameterize these effects.  Second type of problems concerns the degree of approximation used in the actual modeling procedure. A particularly important class of such approximations deals with the description of interaction of radiation and matter in the disk. It should be emphasized that AGN disks, like atmospheres of hot stars, are typical examples of a medium where radiation is not only a {\\em probe} of the physical state, but in fact a crucial {\\em constituent}. In other words, radiation not only carries an information about the medium, it in fact {\\em determines} its structure. Consequently, a treatment of this interaction is in a sense the very gist of the problem. We should therefore study the influence of various approximations, such as the degree of equilibrium assumed, or, more specifically, the extent to which the local thermodynamic equilibrium (LTE) applies, and the completeness of considered opacity and emissivity sources on the computed model, etc.  In the previous paper (Hubeny \\& Hubeny 1997 -- hereafter called Paper I), we have presented some representative self-consistent, non-LTE models of the vertical structure of AGN accretion disks. The basic aim of that study was to investigate the differences in the predicted spectrum between this and previous approaches  (Sun \\& Malkan 1989; Laor \\& Netzer, 1989; Ross, Fabian, \\& Mineshige 1992; Wehrse et al. 1993; St\\\"orzer and Hauschildt 1994; Coleman 1994; Shields \\& Coleman 1994; Blaes \\& Agol 1996; D\\\"orrer et al. 1996). The emphasis was to clarify the role of departures from LTE and to study the effects of simplifications of the hydrostatic equilibrium equation based on assuming a depth-independent vertical gravity acceleration.    In the present paper, we will consider models of vertical structure of AGN disks in detail. In particular, we will study the dependence of computed vertical structure and corresponding emergent spectrum on the adopted value of viscosity and on the degree of sophistication of the modeling procedure. In order to better emphasize the observable consequences of self-consistent, non-LTE models, we present the predicted spectra for a few representative points on the disk. Complete spectra which are obtained by integrating the local spectra over the disk surface, taking into account general relativistic photon transfer functions (Cunningham 1975; Speith, Riffert, \\& Ruder 1995; Agol 1997; Agol, Hubeny, \\& Blaes 1998), will be considered in subsequent papers of this series.  Our aim here is not to construct a model to be used for comparison with actual observations. Instead, we intend to study the sensitivity of computed models on the degree of approximations used in the modeling procedure. Therefore, we will first give a detailed overview of the structural equations and the modeling procedures. This is meant to provide a firm framework on which our approach is based, which in turn is useful to assessing possible systematic effects within our models. ",
        "conclusions": "We have calculated several representative models of vertical structure of an accretion disk around a supermassive Kerr black hole. The interaction of radiation and matter is treated self-consistently, taking into account departures from LTE for calculating both the disk structure and the radiation field. The viscosity is parameterized through the parameter $\\alpha_0$ that describes the vertically averaged viscous stress, and two power--law exponents $\\zeta_0$ and $\\zeta_1$, and the division point $m_{\\rm d}$ between these two forms. The disk structure and emergent radiation is sensitive mainly to the values of $\\alpha_0$, while the other parameters influence the disk structure to a much lesser extent. However, although the detailed shape of the predicted spectrum is sensitive to adopted $\\alpha_0$, the overall appearance of the spectrum is quite similar in all cases.  We have shown that effects of departures from LTE are very important for determining the disk structure and emergent radiation, particularly for hot and electron--scattering dominated disks. We have shown that at least for the disk parameters studied in this paper, NLTE effects typically tend to diminish the value of the Lyman jump; in hot models they suppress the Schuster mechanism by which the LTE models produce a strong emission jump, and in cooler models they increase the flux in the Lyman continuum due to an underpopulation of the hydrogen ground state. Also, we have shown that relaxing the approximation of detailed radiative balance in the hydrogen and helium lines (i.e., computing the so-called NLTE/L models) changes the predicted line profiles significantly, but otherwise does not yield significant changes in computed vertical structure or emergent continuum flux. This result shows that for estimating the {\\em continuum} radiation of AGN disks composed of hydrogen and helium, the NLTE/C models provide a satisfactory approximation.  So far, we have limited our analysis to a simple H-He chemical composition. A preliminary study (Hubeny \\& Hubeny 1998) indicates that the effects of numerous metal lines on the prediced spectral energy distribution of AGN disks may be quite significant. However, that study used a simplified approach in which the vertical structure was fixed by a H-He model, while the line opacity was taken into account only in the spectrum synthesis, assuming LTE source function in metal lines. Such an approach is inconsistent, since, first, the metal lines may change the disk vertical structure (the so-called metal line blanketing effects, long known form the theory of classical stellar atmospheres) and, second, the source function in metal lines may depart significantly from the LTE value. We will therefore need to construct self-consistent, fully metal-line-blanketed models of vertical structure of AGN disks, taking into account effects of literally millions of spectral lines in NLTE. A work of this project is under way, and will be reported in a future paper of this series.  The results presented here do not indicate any fatal flaw of the AGN accretion disk paradigm. In contrast, they show that one of the previous critical arguments against the accretion disk paradigm, the magnitude of the Lyman jump, essentially disappears when increasing a degree of realism of the modeling procedure by relaxing previous simplifying approximations, in particular the local thermodynamic equilibrium and a simplified vertical disk structure. However, this study has concentrated on only one aspect of the problem, the spectral energy distribution in the optical, UV, and EUV region. Many questions, such as the overall spectral energy distribution of the whole disk, the effects of external irradiation, the continuum polarization, etc., remain to be explored in detail. This is exactly what we intend to do in future papers of this series."
    },
    "9804/astro-ph9804241_arXiv.txt": {
        "abstract": "We report on an observation of the low-mass X--ray binary \\U\\ performed during the \\B\\ Science Verification Phase. An absorption feature at $\\sim 37$~keV, attributable to electron cyclotron resonance, has been discovered in its pulse averaged spectrum. The inferred neutron star magnetic field strength is $3.2\\cdot (1+z) \\times 10^{12}$ G, where $z$ is the gravitational redshift. The feature is deep and narrow and is resolved in both the broad-band fit and in the ratio of observed counts to those seen from the Crab. The cyclotron resonance energy is in good agreement with the empirical relation between cyclotron energy and high energy cutoff, while its width is in agreement with the expected Doppler broadening of thermal electrons at the cyclotron resonance frequency. The broad-band 0.1--200 keV spectrum is well fit by a two-component model: a $0.27\\pm 0.02$ keV blackbody and a power law with a photon index of $0.89\\pm 0.02$. This is the first broad-band observation made after the change from spin-up to spin-down that occurred in mid 1990: it confirms the harder spectrum with respect to those observed in the 2--10 keV range. ",
        "introduction": "Following the first detection in Her X--1 \\cite{576}, line features in the 10--110 keV energy range were detected in some X--ray pulsars \\cite{407,375}. These features are interpreted as being due to electron cyclotron transitions in the $\\ga 10^{12}$~G magnetic field of the neutron star \\cite{407}. Because the energy of the fundamental cyclotron harmonic, $E_{\\rm cyc}$, is related to the neutron star magnetic field strength in units of $10^{12}$~G, $B_{12}$, by the relation $E_{\\rm cyc} = 11.6\\, B_{12}\\cdot (1+z)^{-1}$~keV, where $z$ is the gravitational redshift, the observation of cyclotron resonance features (CRFs) in pulsar spectra gives a direct measurement of the magnetic field of the neutron star. From a theoretical point of view CRFs are expected to be visible as absorption features. This was shown {\\em e.g.\\/} by Nagel (1981) \\nocite{306} for any reasonable optical depth. The feature is not strictly due to absorption, but rather to scattering of photons in the wings of the line. Up to now CRFs have been observed mainly in the spectra of young high-mass binary systems, with the notable exception of Her X--1 \\cite{576}.  The 7.7~s X--ray pulsar \\U\\ is one of the few pulsators among LMXRBs. Its timing history is characterized by a sudden change in its spin state that occurred on 1990 June \\cite{1548}. Before this date the source was steadily spinning-up, while from the transition to the present time \\U\\ is spinning-down. Besides coherent pulsation, \\U\\ also shows quasi-periodic oscillations (QPOs) both in X--rays \\cite{29} and in the optical \\cite{1618}, with a centroid frequency of $\\sim 40$ mHz.  The pulse-phase averaged spectrum of \\U\\ has been described in terms of a two component model: an absorbed blackbody with a temperature $kT\\sim 0.6$ keV and a power law \\cite{658}. The spectrum also exhibits a high energy cutoff at $\\sim 20$ keV, interpreted as due to a possible CRF at that energy \\cite{617}, and an emission Ne line complex at $\\sim 1$~keV \\cite{393}. While the spin transition did not affect the high energy part of the spectrum, within the accuracy of previous measurements, it did affect the 2--10 keV part \\cite{1590}: the power law photon index changed from $\\sim 1.6$ \\cite{617} to $\\la 0.7$ \\cite{393}, with a correspondingly significantly lower 2--10 keV flux, that shows an overall fading by a factor four from the HEAO1/A2 measurement \\cite{1548}. ",
        "conclusions": "This is the first \\U\\ broad-band spectrum obtained after the  spin transition. From our flux measurement we can see that the overall {\\em bolometric\\/} flux of the source decreased by a factor 4.1 since the HEAO1/A2 measurements. We find a power law index slightly higher than those found in narrow-band spectra \\cite{393,1612}. This is an effect due to the broad--band fit, in which we simultaneously take into account both the low--energy power law and the high--energy cutoff.  Our \\U\\ CRF measurement, together with its cutoff energy, fits the empirical relation found in X--ray binary pulsars between these parameters: higher the cyclotron resonance energy, higher the cutoff --- and therefore harder the spectrum \\cite{61,1286}. Our CRF measurement also is in agreement with the expected electron Doppler broadening. Indeed, at the cyclotron resonance frequency $\\omega_c$, electrons at rest absorb photons of energy $\\hbar\\omega_c$. For moving, thermal electrons the Doppler broadening $\\Delta\\omega_D$ is predicted to be \\cite{1614} \\begin{equation} \\Delta\\omega_D = \\omega_c \\left( \\frac{2kT_e}{m_ec^2} \\right)^{1/2} |\\cos\\theta| \\label{eq:deltaD} \\end{equation} where $kT_e$ is the electron temperature, and $m_ec^2$ is the electron rest mass. The angle $\\theta$ measures the direction of the magnetic field with respect to the line of sight. Outside the range $\\omega_c\\pm\\Delta\\omega_D$ the cyclotron absorption coefficient decays exponentially, and other radiative processes become important. From Eq.~\\ref{eq:deltaD} and the CRF parameters given in Table~\\ref{tab:fit} we obtain a lower limit to the electron temperature of $\\sim 5$ keV, in reasonable agreement with the calculations of self-emitting atmospheres of Harding et~al.\\ (1984) \\nocite{304}, according to which the temperature for a column or slab with optical depth $\\sim 50$ g~cm$^{-2}$ is 4--8$\\times 10^7~^\\circ$K, corresponding to $kT_e\\sim 3.5$--7 keV.  Finally, we want to mention that if we assume the QPO frequency as due to the beating between the pulse frequency and the Keplerian motion at the magnetospheric radius, then the \\U\\ magnetic field strength is related to its luminosity by the relation $B_{12}\\sim 5.5\\sqrt{L_{37}}$ \\cite{29}, where $L_{37}$ is the X--ray luminosity in units of $10^{37}$~erg~s$^{-1}$. Assuming a source distance $5\\la \\rm d_{\\rm kpc} \\la 13$ \\cite{1618}, this corresponds to $2.4\\la B_{12} \\la 6.3$, in  good agreement with our measurement."
    },
    "9804/gr-qc9804024_arXiv.txt": {
        "abstract": "\\noindent In the coming decade, the LIGO/VIRGO network of ground-based kilometer-scale laser interferometer gravitational wave detectors will open up a new Astronomical window on the Universe: gravitational waves in the frequency band $10$ to $10^4$ Hz.  In addition, if the proposed, 5 million kilometer long, space based interferometer LISA flies, another window will be opened in the frequency band $10^{-4}$ to $1$ Hz.  We review the various possible sources that might be detected in these frequency bands, and the information that might be obtainable from observed sources.  Several key possible sources are inspirals and coalescences of neutron-star neutron-star and/or neutron-star black-hole binaries; inspirals, mergers, and ringdowns of black-hole black-hole binaries (both solar mass and supermassive); stellar core collapse; rapidly rotating neutron stars; the formation of supermassive black holes; and inspirals of compact objects into supermassive black holes. ",
        "introduction": "This review of gravitational wave sources is divided into three sections.  First, we review the detector sensitivities that have been achieved to date and discuss projected sensitivities for detectors now under construction.  Second, we summarize the current observational upper limits on gravitational waves in various frequency bands.  The main body of this review will consist of a survey of various anticipated sources of waves. Each anticipated source can be roughly characterized by a characteristic frequency $f$ and a characteristic value of strain amplitude $h$.  However, it is important to also note that sources vary widely with respect to how uncertain is the rate of their occurrence in the Universe.  The enterprise of anticipating potential gravitational wave sources is extremely uncertain.  For most sources that we can conceive of, either the wave strengths are uncertain by several orders of magnitude, or the event rate is uncertain by several orders of magnitude, or the very existence of the source itself is very uncertain.  While there are some important exceptions such as coalescing compact binaries, these are the exception rather than the rule.  The upside of this great uncertainty is the potential for gravitational wave  astronomy to bring us new and interesting information.     For more details on the topics discussed here, the reader is encouraged to consult the detailed recent review articles by Thorne \\cite{300yrs,thorne95,thorne97a,thorne97b}, and also the review article by Allen \\cite{allen96} on stochastic gravitational waves. ",
        "conclusions": "With the kilometer scale interferometer network about to come online in the next few years, this is an exciting time for gravitational wave astronomy.  Almost certainly we have not anticipated all the detectable sources of gravitational waves that Nature produces in the real Universe.  Hopefully she will bring us some surprises."
    },
    "9804/astro-ph9804063_arXiv.txt": {
        "abstract": "We analyze the linear, 3D response to tidal forcing of a disk that is thin and thermally stratified in the direction normal to the disk plane.  We model the vertical disk structure locally as a polytrope which represents a disk of high optical depth.  We solve the 3D gas-dynamic equations semi-analytically in the neighborhood of a Lindblad resonance.  These solutions match asymptotically on to those valid away from resonances (previously obtained by Korycansky \\& Pringle 1995) and provide solutions valid at all radii $r$.  We obtain the following results.  1) A variety of waves are launched at resonance, including r modes and g modes. However, the f mode carries more than 95\\% of the torque exerted at the resonance.  2) These 3D waves collectively transport exactly the amount of angular momentum predicted by the standard 2D resonant torque formula.  3) Near resonance, the f mode behaves compressibly and occupies the full vertical extent of the disk. Away from resonance, the f mode behaves incompressibly, becomes confined near the surface of the disk, and, in the absence of other dissipation mechanisms, damps via shocks.  In general, the radial length scale for this process is roughly $r_{\\rm L}/m$ (for resonant radius $r_{\\rm L}$ and azimuthal tidal forcing wavenumber $m$), {\\it independent of the disk thickness $H$}. This wave channeling process is due to the variations of physical quantities in $r$ and is not due to wave refraction.  4) However, the inwardly propagating f mode launched from an $m=2$ inner Lindblad resonance experiences relatively minor channeling (accompanied by about a factor of 5 increase in nonlinearity), all the way to the radial center of the disk.  We conclude that for binary stars, tidally generated waves at Lindblad resonances in highly optically thick circumbinary disks are subject to strong nonlinear damping by the channeling mechanism, while those in circumstellar accretion disks are subject to weaker nonlinear effects.  We also apply our results to waves excited by young planets for which $m \\approx r/H$ and conclude that the waves are damped on the scale of a few $H$. ",
        "introduction": "Gaseous disks are found in many types of binary star systems, including cataclysmic variables (CVs) and pre-main-sequence stars, and young planetary systems.  The orbiting objects (stars or planets) exert tidal forces on these disks, which generally act merely to distort the disks from an axisymmetric form. However, at special locations in disks where resonances occur, the tidal forces generate waves that transport energy and angular momentum.  As a result, a resonant torque is exerted by the system objects.  The orbital evolution of the perturbing objects sometimes depends on the strength of these torques (Goldreich \\& Tremaine 1980, hereafter GT80; Lin \\& Papaloizou 1993; Lubow \\& Artymowicz 1996).  The waves transfer their angular momentum and energy to the disk in the regions of space where they damp, and this in turn affects the evolution of the disk.  For example, gaps could be created in disks in regions where the waves damp.  This paper concentrates on Lindblad resonances (LRs), which are due to horizontal forcing (along the disk plane).  We assume throughout that the disk is coplanar with the orbit of the system objects.  A 2D, linear theory for resonant tidal torques and associated wave propagation was developed by Goldreich \\& Tremaine (1979, hereafter GT79). This theory provides an explicit formula for the Lindblad resonant torque.  The 2D theory considers the disk to have only radial and azimuthal extent and ignores effects over its vertical extent (perpendicular to the orbit plane).  Important progress has been made in the study of 2D nonlinear waves in disks (Shu, Yuan, \\& Lissauer 1985; Yuan \\& Cassen 1994). The torque in the nonlinear case was found to be within a few percent of that predicted by the 2D linear formula. The nonlinearities produce highly spiked density profiles which increase the level of dissipation present in a viscous disk. Radiative damping of linear waves can be important particularly when the disk is warm (Cassen \\& Woolum 1996). Turbulent viscosity in the disk provides another means of wave damping.  The 2D treatment is valid if both the vertical structure of the disk and its thermodynamic response are locally isothermal.  Under such circumstances, tidal forcing will generate a 2D wave in a 3D disk. The wave front remains perpendicular to the disk plane at all heights, as the wave propagates radially in the disk.  However, this 2D wave is highly singular in that it does not exist in a disk with a vertical temperature variation (Lin, Papaloizou, \\& Savonije 1990a, hereafter LPS; Lubow \\& Pringle 1993, hereafter LP).  Furthermore, a vertically isothermal structure is not realistic for many important classes of gaseous disks, such as accretion disks in CVs, circumstellar and circumbinary disks of YSOs, and protoplanetary disks. Such disks often have optical depths much greater than unity and can be expected to have substantial vertical temperature variations, if they have an internal heat source such as turbulent dissipation.  We demonstrate in this paper that 2D tidal forcing, caused by LRs, excites 3D waves in a thermally stratified disk.  However, another class of resonances exist due to 3D effects. These resonances are due to vertical, tidal forcing by a coplanar perturber of a disk with nonzero thickness (Lubow 1981).  The vertical resonances also generate horizontally propagating waves. Although intrinsically weaker than the LRs, the vertical resonances may be of importance in close binary star systems.  Some investigations of 3D effects have been carried out using numerical simulations (Lin, Papaloizou, \\& Savonije 1990a,b).  Such approaches can be used to explore a limited range of physical parameter space. However, recent progress has been made in obtaining semi-analytic solutions for waves in 3D disks (LP; Korycansky \\& Pringle 1995, hereafter KP). The aim of this paper is to extend that approach to understand 3D wave generation at LRs and the subsequent wave propagation.  We determine the linear response of a thin (but nonzero thickness) disk.  We model the disk locally as a polytrope in the vertical direction, which is valid for a disk of very high optical depth, and we ignore the effects of atmospheric layers.  We aim to understand which modes of a thermally stratified disk are excited at LRs and how much torque is carried by such waves. By studying the properties of linear wave propagation, we can also understand where nonlinearity sets in that will likely lead to shocks and subsequent wave dissipation.  The outline of this paper is as follows. In \\S2, we review the properties of disk modes. In \\S3 and \\S4, we derive and solve the equations for waves generated at LRs.  In \\S5, we compute the total torque carried by these waves and determine which modes are excited. In \\S6, we summarize properties of the f mode of the disk and present our numerical results.  In \\S7, we discuss the application of our results to binary and protoplanetary systems. \\S8 contains a summary. ",
        "conclusions": "\\subsection{Summary of results}  We have analyzed the 3D response to tidal forcing of a gaseous disk having a vertical temperature variation. The unperturbed vertical disk structure was modeled as locally polytropic.  We have considered effects at Lindblad resonances (LRs), subject to the thin disk approximation $H \\ll r/m$, for which the LRs provide horizontal forcing that is independent of height.  In the standard vertically isothermal case considered to date, the horizontally propagating 2D mode (sound wave) carries all the resonantly generated angular momentum, but only if the thermodynamic response of the disk is also isothermal.  In the vertically polytropic case, the f mode has a role equivalent to that of the 2D mode, although r and g modes are also launched at LRs (see Figure~1). The mode launched at vertical resonances (Lubow 1981) can be shown to be the ${\\rm p}_1^{\\rm e}$ mode.  We have found that nearly all the torque exerted at an LR is carried by the f mode (see Table~1). Near resonance, the f mode behaves in a similar manner to the 2D mode in the isothermal case, in that it occupies the full vertical extent of the disk and behaves somewhat compressibly.  The f mode is almost two-dimensional near resonance in that the vertical velocities are smaller than the horizontal by a factor of order $(H/r)^{1/3}$.  However, away from resonance, the behaviors of the f and 2D modes differ radically (see \\S6.2).  At a distance $\\sim r_{\\rm L}/m$ from the resonance, the f mode begins to concentrate its energy near the surface of the disk, a process which we have called wave channeling. In this regime, the f mode behaves like a surface gravity mode (see Plates~1 and~2). In most cases, the wave amplitude increases by many orders of magnitude with  distance from resonance, so that it is very likely that shocks will develop which would damp the wave.  An important exception to this last statement occurs for the $m=2$ ILR. The wave generated there undergoes relatively mild wave channeling and consequently a relatively mild increase in nonlinearity (by about a factor of 5). The amount of increase depends on the density distribution in the disk and could even result in a decrease in dimensionless amplitude near the disk center. Dissipation by turbulent viscosity is likely important in some cases, such as CV disks. On the other hand, the wave generated at the $m=2$ OLR does exhibit strong wave channeling.  These results indicate that circumbinary disks, as found around young binaries, and protoplanetary disks perturbed by planets, are subject to strong effects of wave channeling (see \\S7).  \\subsection{Discussion}   The results in this paper differ from the previously accepted picture for wave propagation in thermally stratified disks (e.g. LPS). The standard expectation was that a wave launched at a Lindblad resonance would begin propagating horizontally, but the wavefront would be rapidly refracted upwards as a result of the decrease in sound speed with increasing height above the mid-plane.  After advancing a distance comparable to the disk thickness $H$, the wavefront would be substantially tilted upwards.  The wave would then propagate vertically into the atmosphere of the disk, where it would shock. This model is based on the idea that the launched wave is a pressure wave or p mode in a high-frequency acoustic limit, so that one can consider the wavefront to propagate at the local sound speed without being affected by inertial or buoyancy forces.   Our results provide a different picture.  We have found that the wave cannot be considered to propagate vertically, since it is in fact a vertically evanescent f mode. In this view, taken in LP and KP, the disk behaves like a waveguide in which the wave is vertically confined. However, somewhat similar to the standard picture, the wave energy does rise to the surface of the disk as the wave propagates away from the resonance.  But this wave channeling process is effective over a distance of order $r_{\\rm L}/m$, where $r_{\\rm L}$ is the radius of the Lindblad resonance and $m$ is the azimuthal wavenumber.  It does not depend on the disk thickness, provided that $H/r \\ll1$.  This can be understood roughly as follows.  The wave channeling occurs because of the radial variation of the dimensionless intrinsic frequency of the wave, $(\\omega-m\\Omega)/\\Omega$.  This is equal to $\\pm\\kappa/\\Omega$ at the Lindblad resonance, but increases rapidly in magnitude over a distance of order $r_{\\rm L}/m$.  The wave therefore proceeds rapidly along the f-mode branch of the dispersion relation. In the high-frequency limit, as described in \\S6.1, the mode behaves like a surface gravity wave and is confined near the surface of the disk.  (This would be true even in an incompressible disk in which refraction cannot operate.)  The surface-gravity-mode dispersion relation is approximately $(\\omega -m \\Omega)^2 \\approx g_{\\rm s} k$, where the surface gravity $g_{\\rm s}=\\Omega^2 H$ and $k$ is the radial wavenumber.  It can be seen (see \\S6.1) that the mode is confined to the disk surface in a layer of vertical thickness $\\delta  \\sim k^{-1}$.  It then follows from the dispersion relation that $(\\omega-m\\Omega)^2/\\Omega^2 \\sim H/\\delta $.  As a result, the wave becomes confined (channeled) near the disk surface (i.e., $\\delta /H \\la 1/2$) over a radial distance from resonance of order $r_{\\rm L}/m$.  Another issue concerns the boundary conditions.  In LPS, the anticipated shocks in the disk atmosphere were represented in their simulation by an upper boundary condition that included some amount of dissipation.  In our analysis of a polytropic disk, we have applied a boundary condition at the disk surface that acts to reflect waves rather than absorb energy. Even in a purely isothermal disk without a definite surface, as demonstrated by LP, all the modes are confined vertically and do not propagate vertically to infinity.  This is a consequence of the increase in vertical gravity with height.  Our treatment of the boundary is valid in the limit of high optical depth, where the atmosphere occupies a negligible mass.  In particular, the effects of atmosphere are not strongly felt by the wave until $k H_{\\rm atmos} > 1$, with $H_{\\rm atmos}$ being the density scale height at the base of the atmosphere. For a highly optically thick disk, $H$ is much greater than $H_{\\rm atmos}$, and so $k H \\gg 1$. When this condition occurs, the wave has been strongly channeled before much wave energy enters the atmosphere. We plan to extend our current analysis to include the effects of a disk atmosphere in a future paper.    The asymmetry in wave properties between the inner and outer LRs is most pronounced for low-$m$ cases. For fixed $m$, the level of nonlinearity for ILR waves is lower than for OLR waves (see \\S6.2.2). One reason is that the density increases as the ILR wave propagates inward, while the density decreases as the OLR wave propagates outwards (see eq. [\\ref{calN2}]). The other cause of the asymmetry is that the f-mode dispersion relation is followed to higher wavenumbers in the exterior of the OLR than in the interior of the ILR, for low $m$ (see Figure~3). Away from the LR, the f mode is vertically confined to a surface layer of thickness $\\sim k^{-1}$. Due to wave action conservation, the confinement acts to increase the wave amplitude, causing substantial asymmetry in the case $m=2$.  In the high-$m$ case, the behavior of the wavenumber becomes similar near the inner and outer LRs, and consequently the nonlinearity becomes important at comparable distances from the respective resonances.   A variant on the vertically polytropic model occurs for layered disks (Gammie 1996). In this model, some regions of protostellar or protoplanetary disks have turbulence restricted to the upper layers of the disk. The resulting vertical temperature structure is isothermal in the non-turbulent layer near mid-plane and the temperature declines with height in the outer turbulent layers. A wave launched from an LR in such a disk would probably require a longer distance from resonance to undergo wave channeling.  Given the large increases in wave velocities associated with wave channeling (see Plates~1 and~2), we expect that nonlinearities will cause the waves to shock and deposit angular momentum preferentially near the surface of the disk. We have described where nonlinearities are expected to become important in \\S6.2, based on a quasilinear estimate. Little work has been carried out to investigate the nonlinear behavior of f modes or surface gravity waves in a compressible fluid. It would be useful to conduct nonlinear simulations of the waves in such a disk."
    },
    "9804/astro-ph9804313_arXiv.txt": {
        "abstract": "We have used the Deep Survey telescope of EUVE to investigate shadows in the diffuse EUV/Soft X-Ray background cast by clouds in the interstellar medium.  We confirm the existence of a shadow previously reported, and provide evidence for two new shadows. We used IRAS data to identify the clouds producing these shadows and to determine their optical depth to EUV radiation. The EUV-absorbing clouds are optically thick in the EUV, and all EUV emission detected in the direction of these shadows must be produced from material in front of the clouds.  We obtained new optical data to determine the distance to these clouds. We use a new differential cloud technique to obtain the pressure of the interstellar medium. These results do not depend on any zero level calibration of the data. Our results provide evidence that the pressure of the hot interstellar gas is the same in three different directions in the local interstellar medium, and is at least 8 times higher than derived for the local cloud surrounding our Sun. This provides new evidence for large thermal pressure imbalances in the local ISM, and directly contradicts the basic assumption of thermal pressure equilibrium used in almost all present models of the interstellar medium. ",
        "introduction": "The discovery of the soft X-ray / EUV background (\\cite{bow68}) and its anti-correlation with the Galactic H{\\sc i} distribution has led to an on-going debate on the origin of that emission. Since an absorption column density of $N_{\\rm H{\\sc i}} = 1.2 \\times 10^{20}$\\,cm$^{-2}$ (which corresponds to an optical depth $\\tau = 1$ at 0.25\\,keV) is reached at a distance of approximately 100\\,pc in the direction of the Galactic plane, the diffuse background must arise from hot gas in the local interstellar medium (ISM). Optical and UV absorption line measurements of nearby stars together with observations in the X-ray and EUV wavelength range have established a local region with a significant H{\\sc i} deficiency which is partly filled with a plasma, the so-called ``Local Bubble'' (LB). No definite model exists for this region and its origin is unclear.  One view is that the gas in the LB has been heated by a supernova explosion which occurred in the solar neighborhood about $10^7$\\,yr ago (cf. \\cite{mckost77}, \\cite{coxrey87}). Other authors claim that the LB is part of an asymmetrically shaped superbubble created by stellar winds and supernova explosions (cf. \\cite{frisch95}). However, the physical properties such as density, temperature, pressure, and extension of the LB, are not well determined and so far observational constraints are insufficient to establish a canonical model of the evolution and the origin of the LB. The McKee \\& Ostriker model is the best known model but it is not easily testable. Only soft X-ray and EUV observations provide a direct method to study the hot gas in the ISM and test LB models.  This paper presents new measurements for the thermal pressure in the local ISM, using observations of nearby clouds taken by the Deep Survey (DS) telescope on-board the {\\em Extreme Ultraviolet Explorer} (EUVE). Prior to this paper, Bowyer~et~al.~(\\cite{bow95}) reported on the first detection of a spatial absorption feature in the diffuse EUV background; this feature was positionally coincident with an IRAS cirrus cloud at a distance of $\\la$ 40\\,pc. Since the cloud casting this shadow is optically thick at EUV wavelengths, the authors concluded that all residual EUV emission observed at the position of the cloud originated from material along the line of sight in front of the cloud. By attributing the background subtracted flux in front of the cloud to the hot gas in the ISM, Bowyer~et~al.~(\\cite{bow95}) derived for the length of the emitting region (40\\,pc) a pressure (P/$k$) of 19,000\\,cm$^{-3}$K. These authors pointed out that an imperfect zero level calibration of the data can lead to a smaller value for the pressure, and estimated a lower limit for the pressure of 7,000\\,cm$^{-3}$K.  The purpose of this paper is to present a re-observation of the previously reported cloud shadow in the EUV background and observations of two new shadows discovered in the DS. For all three shadow regions we provide new optical photometry data which establishes the distance of the respective clouds. We employ a differential cloud technique which does not depend on any zero level calibration, to obtain the pressure of the ISM in these directions.  The outline of this paper is as follows: First, in Sect.~2 we describe our new observations and the data analysis. In Sect.~3, we employ a new analysis technique and compare our results for the three cloud shadows. We derive physical quantities of the hot interstellar gas from our observations. Finally, in Sect.~4, we discuss the implications for models of the local ISM. ",
        "conclusions": "\\label{discus}  We have presented a re-observation of a cloud shadow in the diffuse EUV background discovered by Bowyer~et~al.~(\\cite{bow95}) and present two additional shadows detected in the EUVE DS.  We have developed a new method to derive the diffuse astrophysical EUV background from the EUVE DS data. This differential cloud technique can be applied to any cloud shadow data without a known zero level calibration. In our case we can extrapolate the results to the origin and obtain a zero level calibration for the DS detector. This zero level calibration indicates that about 50\\% of the detected photons are astronomical in origin. More importantly, this zero level calibration allows a determination of the diffuse EUV background in the direction of the three nearby cloud shadows and hence to derive the pressure in the regions between us and these clouds. The pressure obtained for these nearby regions is consistent with the value obtained from the differential measurements for the more distant regions along the lines of sight.  The results of our cloud shadow observations with EUVE provide evidence for a constant pressure in three different directions in the local ISM. We derived a pressure of P/$k \\approx 16500 {\\rm cm}^{-3}$\\,K. The canonical value obtained from the analysis of solar He~{\\sc i} 584\\,\\AA~~radiation resonantly scattered by helium in the inflowing cloud is P/$k = 730 \\pm 30 {\\rm cm}^{-3}$\\,K (e.g., \\cite{frisch95}). Based on line measurements with EUVE, Vallerga (\\cite{vallerga96}) derived a pressure for the solar cloud in the range P/$k = 1700 - 2300 {\\rm cm}^{-3}$\\,K which is in good agreement with results obtained by \\cite{bertetal85}, P/$k \\approx 2600 {\\rm cm}^{-3}$\\,K). A comparison of these results shows the pressure of the local ISM exceeds the pressure of the cloud surrounding our Sun by a factor of $\\ge 8$. The original observation (Bowyer~et~al.~\\cite{bow95}) of a pressure imbalance between the solar cloud and the surrounding local ISM could have been the result of a nearby shock wave which had not yet reached the Sun. The results reported here are obtained in three different directions. The isotropy in the results for all these directions is not consistent with this hypothesis.  Flux in the soft X-ray bandpass requires a larger absorption column density to reach unit optical depth and it is difficult or impossible to separate the locally produced diffuse soft X-ray emission from contributions originating at larger distances. Soft X-ray observations, therefore, typically provide only an upper limit for the pressure of the local ISM. As discussed previously, the pressure obtained from our cloud shadow observations in the EUV is strictly interpreted as a lower limit. Although different plasma codes have been used by different authors, the pressure derived from our cloud shadow observations with EUVE is about the same as values derived from soft X-ray observations (e.g., \\cite{snowden97}, \\cite{freyberg97})when we use the same code (P/$k$(EUV) = 13500 cm$^{-3}$K; P/$k$(soft X-ray) = 14000 cm$^{-3}$K). The near equality of these upper and lower limits implies we have measured the true pressure.  Beyond the cloud surrounding our Sun we found no significant intrinsic cold gas absorption in the local ISM. This supports the idea of a cavity primarily filled with an ionized plasma. Interstellar absorption lines in the spectra of nearby stars have indicated the existence of cloudlets of cold gas (like the local cloud around our Sun) in the solar vicinity (Lallement~\\cite{lallement96}). However, the effect of these cloudlets on the diffuse EUV emission observed in the directions of the cloud shadows discussed here is negligible. We point out that the different line of sights for the six clouds discussed here show that the local ISM cannot be simply described as a sphere with an almost constant radius (\\cite{snowcox90}).  Our observations provide evidence for a large pressure imbalance in the local ISM compared to the local cloud surrounding our Sun. This contradicts the basic assumption of almost all available models for the local ISM of a pressure equilibrium (e.g., \\cite{mckost77}, \\cite{coxrey87}). It is unclear what physical mechanism can maintain such a large pressure imbalance. We emphasize that these models as well as the calculations in the here presented work assume a hot plasma in collisional ionization equilibrium. A completely different approach is provided by Breitschwerdt \\& Schmutzler (\\cite{breischm94}). Their model calculations are based on an adiabatically cooling non-equilibrium plasma. The advantage of this model is that it can at least qualitatively explain many of the observed features of the local ISM and it does not require pressure equilibrium in the local ISM."
    },
    "9804/astro-ph9804125_arXiv.txt": {
        "abstract": "It has been recently realized (Rephaeli 1995) that the relativistic corrections to the spectral distortions of the cosmic microwave background (CMB) measured in the direction of clusters of galaxies containing hot gas are significant and should be detectable with the forthcoming experiments. In the present paper we calculate the correction terms that are proportional to $V_{\\rm r}/c\\times kT_e/m_ec^2$ and $(V/c)^2$ to the standard formulae describing the spectral distortions caused by the bulk motion of the free electrons (kinematic effect) and due to the presence of the hot gas (thermal effect) for the case of a cluster having a peculiar velocity $V$ ($V_{\\rm r}$ is its radial component). The results of our analytical calculations are confirmed by Monte-Carlo simulations (Sazonov \\& Sunyaev 1998). ",
        "introduction": "Thomson scattering of the cosmic microwave background (CMB) radiation by hot electrons in the intergalactic gas in clusters of galaxies modifies the spectrum of the CMB (Sunyaev and Zel'dovich 1972). Zel'dovich and Sunyaev (1969), basing on the Kompaneets equation (1957), derived a simple formula describing the spectral form of the distortion, which is proportional to the parameter $y=(kT_e/m_ec^2) \\tau$, where $\\tau$ is the Thomson optical depth along the line of sight. The effect has now been observed from a number of clusters of galaxies (see Birkinshaw 1998 for review).  Recently, interest to this effect has been reactivated in view of the perspectives of accurate measurement of the CMB distortions in a number of experiments, both ground-based and on balloons, by the MAP spacecraft and especially by the Planck Surveyor mission scheduled to be flown in the middle of the next decade. These activities were motivated by the fact that the gas temperature is so high in the clusters of galaxies (ranging between 3 and 17~keV, Tucker et al. 1998) that the scattering electrons have thermal velocities of the order of 0.1 -- 0.3 $c$, so one has to include into consideration the relativistic corrections to obtain an accurate result. Rephaeli (1995), basing on extensive previous work (Wright 1979, Fabbri 1981, Taylor \\& Wright 1989, Loeb et al. 1991) has demonstrated by means of numerical calculations the relevance of the relativistic corrections for the future experiments. Stebbins (1997), Itoh et al. (1998), and Challinor \\& Lasenby (1998) used a Fokker-Planck approximation of the relativistic photon kinetic equation to obtain corrections, written as series in powers of $kT_e/m_ec^2$, to the standard nonrelativistic solution. These results have proved to be in excellent agreement with those of Rephaeli, demonstrating the applicability of the diffusion approximation to the problem at hand, despite the small optical depths of the clusters of galaxies ($\\tau \\sim 0.01$).  A gas cloud moving rapidly relative to the CMB along the observer's line of sight must significantly modify the spectrum of the CMB in addition to the thermal effect. The change in the brightness temperature caused by this ``kinematic'' effect is to first order simply proportional to the radial component of the cluster velocity $\\sim (V_{\\rm r}/c)\\tau$ (Sunyaev \\& Zel'dovich 1980). The effect should be detectable in the future, and will enable measurement of cluster peculiar velocities, with significant implications for studies of the large-scale structure of the universe. It is obvious that corrections similar to those found for the thermal effect must exist and should be taken into account for the kinematic effect, if one wants to find the correct solution for the case of a moving cluster. In this paper we calculate the next-order changes in the spectrum of the CMB related to the cluster peculiar velocity by solving the photon kinetic equation. We have obtained simple formulae giving the correction terms of the orders of $V_{\\rm r}/c\\times kT_e/m_ec^2$ and $(V_{\\rm r}/c)^2$. Our method is similar to that used by Psaltis \\& Lamb (1997) who considered the more general problem of comptonization in a moving media. The solution these authors have obtained, although applicable to many astrophysical situations, does not contain the $O(V_{\\rm r}/c\\times kT_e/m_ec^2)$ term, because this term is third-order in electron velocity, whereas their solution is accurate only to second order in it. We aslo confirm the existence of the term of order $(kT_e/m_ec^2)^2$ found earlier using techniques different from ours (Rephaeli 1995, Stebbins 1997, Itoh et al. 1998, and Challinor \\& Lasenby 1998). Earlier we have found all the correction terms mentioned above numerically using Monte-Carlo simulations (Sazonov \\& Sunyaev 1998). ",
        "conclusions": ""
    },
    "9804/astro-ph9804196_arXiv.txt": {
        "abstract": "We present a new calculation of the intergalactic \\gray\\ pair-production absorption coefficient as a function of both energy and redshift. In reexamining this problem, we make use of a new {\\it empirially} based calculation (as opposed to previous {\\it model} calculations) of the intergalactic infrared radiation field (IIRF). We find smaller opacities than those given previously (Stecker \\& De Jager 1997). We apply our results to the new observations of the flaring \\gray\\ spectra of Mrk 421 and Mrk 501, both at a redshift of $\\sim$0.03. Our new calculations indicate that there should be no significant curvature in the spectra of these sources for energies below 10 TeV, as indicated by recent observations. However, the intrinsic spectra of these sources should be harder by amounts of $\\sim$ 0.25 to 0.45 in the spectral index (in the 1 - 10 TeV range), with an intergalactic absorption cutoff above $\\sim 20$ TeV. ",
        "introduction": "We have previously pointed out (Stecker, De Jager \\& Salamon 1992 (hereafter, SDS92)) that very high energy \\gray\\ beams from blazars can be used to measure the intergalactic infrared radiation field, since pair-production interactions of \\grays\\ with intergalactic IR photons will attenuate the high-energy ends of blazar spectra. Determining the intergalactic IR field, in turn, allows us to model the evolution of the galaxies which produce it. As energy thresholds are lowered in both existing and planned ground-based air Cherenkov light detectors (Cresti 1996), cutoffs in the \\gray\\ spectra of more distant blazars are expected, owing to extinction by the IIRF. These can be used to explore the redshift dependence of the IIRF (Stecker \\& Salamon 1997; Salamon \\& Stecker 1998). Furthermore, by using blazars for a determination of attenuation as a function of redshift,  combined with a direct observation of the IR background from the {\\it DIRBE}  detector on {\\it COBE}, one can, in principle, measure of the Hubble constant $H_{0}$ at truly cosmological distances (Salamon, Stecker \\& De Jager 1994).  There are now over 50 grazars which have been detected by the {\\it EGRET} team (Thompson, \\etal\\ 1996). These sources, optically violent variable quasars and BL Lac objects, have been detected out to a redshift greater that 2. Of all of the blazars detected by {\\it EGRET}, only the low-redshift BL Lac, Mrk 421, has been seen by the Whipple telescope. The fact that the Whipple team did not detect the much brighter {\\it EGRET} source, 3C279, at TeV energies (Vacanti, \\etal\\ 1990, Kerrick, \\etal\\ 1993) is consistent with the predictions of a cutoff for a source at its much higher redshift of 0.54 (see SDS92). So too is the recent observation of two other very close BL Lacs ($z < 0.05$), {\\it viz.}, Mrk 501 (Quinn, \\etal\\ 1996) and 1ES2344+514 (Catanese, \\etal\\ 1997) which were too faint at GeV energies to be seen by {\\it EGRET}.  In this paper, we calculate the absorption coefficient of intergalactic space using a new, empirically based calculation of the spectral energy distribution (SED) of intergalactic low energy photons (Malkan \\& Stecker 1998; hereafter MS98) obtained by integrating luminosity dependent infrared spectra of galaxies over their luminosity and redshift distributions. After giving our results on the \\gray\\ optical depth as a function of energy and redshift out to a redshift of 0.3, we apply our calculations by comparing our results with recent spectral data on Mrk 421 as reported by McEnery, \\etal\\ (1997) and spectral data on Mrk 501 given by Aharonian, \\etal\\ (1997). The results presented here supercede those of our previous calculations (Stecker \\& De Jager 1997), which were based more on theoretical models (see discussion in MS98). We consider the results presented here to be considerably more reliable than any presented previously. ",
        "conclusions": "We have calculated the absorption coefficient of intergalactic space from pair-production interactions with low energy photons of the IIRF, both as a function of energy and redshift, using new, more reliable estimates of the SED for the IIRF which were obtained by MS98. Our results predict less absorption than we obtained previously (Stecker \\& De Jager 1997), because the MS98 IIRF is lower than our previous estimate. One reason for this difference is that our previous IIRF spectrum was normalized partly to reflect the estimate of Gregorich, {\\it et al.} (1995) of the IIRF at 60 \\mic. According to Bertin, Dennefeld \\& Moshir (1997), that estimate may have been based on the inclusion of false detections in their analysis of the IRAS data. For absorption calculations, it is important to use the most reliable estimate of the IIRF avaliable, since the absorption effect depends exponentially on the magnitude of the IIRF. While we do not claim that our new results for the absorption coeeficient as a function of energy differ dramatically from those obtained previously (MacMinn \\& Primack 1996; Stecker \\& De Jager 1997), we {\\it do} claim that they are more reliable because they are based on the empirically derived IIRF given by MS98, whereas all previous calculations of TeV $\\gamma$-ray absorption were based on theoretical modeling of the IIRF. The MS98 calculation was based on data from nearly 3000 IRAS galaxies. These data included (1) the luminosity dependent infrared SEDs of galaxies, (2) the 60$\\mu$m luminosity function of galaxies and, (3) the redshift distribution of galaxies.  We have applied our absorption calculations to recent flaring spectra of the nearby BL Lac objects Mrk 421 and Mrk 501. The spectral calculations given here are in good agreement with the recent observations that indicate no significant curvature in the spectra of Mrk 421 and Mrk 501. The observations are also consistent with our calculated steepening of 0.25 to 0.45 in the spectral index of these sources in the 1-10 TeV range. Our new calculations predict a significant intergalactic absorption effect which should cut off the spectra of Mrk 421 and Mrk 501 at energies greater than $\\sim$20 TeV. Observations of these objects at large zenith angles, which give large effective threshold energies, may thus demonstrate the effect of intergalactic absorption.  Our new calculations confirm the conclusion in SDS92 that TeV spectra of sources at redshifts higher than 0.1 should suffer significant absorption. The recent detection of another XBL at a redshift below 0.1, {\\it viz.}, 1ES2344+514 (Catanese, \\etal\\ 1997), further supports the argument that nearby XBLs may be the only significant TeV sources presently detectable (Stecker, De Jager \\& Salamon 1996)."
    },
    "9804/astro-ph9804019_arXiv.txt": {
        "abstract": "The Hubble Deep Field South (HDFS) has been recently selected and the observations are planned for October 1998.  We present a high resolution (FWHM $\\simeq 14$ \\kms) spectrum of the quasar J2233--606 ($z_{em}\\simeq2.22$) which is located 5.1 arcmin East of the HDFS. The spectrum obtained with the New Technology Telescope redward of the Lyman--$\\alpha$ emission line covers the spectral range 4386--8270 \\AA. This range corresponds to redshift intervals for CIV and MgII intervening systems of $z=1.83-2.25$ and $z=0.57-1.95$ respectively. The data reveal the presence of two complex intervening CIV systems at redshift $z=1.869$ and $z=1.943$ and two complex associated ($z_{abs} \\approx z_{em}$) systems. Other two CIV systems at $z=1.7865$ and $z=2.077$, suggested by the presence of strong Lyman--$\\alpha$ lines in low resolution ground based and Hubble Space Telescope (HST) STIS observations (Sealey et al. 1998) have been identified. The system at $z=1.943$ is also responsible for the Lyman limit absorption seen in the HST/STIS spectrum.  The main goal of the present work is to provide astronomers interested in the Hubble Deep Field South program with information related to absorbing structures at high redshift, which are distributed along the nearby QSO line of sight. For this purpose, the reduced spectrum, obtained from three hours of integration time, has been released to the astronomical community. ",
        "introduction": "The Hubble Deep Field (HDF) program (Williams et al. 1996), which was carried out with the Hubble Space Telescope (HST) in December 1995 in the four filters $UBVI$ to the deepest ever reached limiting magnitude, together with followup observations in other spectral bands, can be considered as one of the major astronomical events of the nineties.  More than 40 articles on this program have been published in refereed journals between April 1996 and December 1997, confirming the tremendous impact of the observations. Articles related to the high redshift Universe ($z\\gsim1$) treat such subjects as the galaxy redshift distribution (Gwyn \\& Hartwick 1996, Seidel et al. 1996, Lanzetta et al. 1996, Lowenthal et al. 1997), the Global Star Formation History (Madau et al. 1996, Connoly et al. 1997, Guzman et al. 1997) and the clustering properties of galaxies (Colley et al. 1996, Villumsen et al. 1997).  After two years, the Hubble Deep Field South (HDFS) has been planned for Cycle 7 of the HST (Williams et al. 1997, see also the HDF--South Web site at the URL \\\\ http://www.stsci.edu/ftp/science/hdf/hdfsouth/hdfs.html), and the observations will be performed in October 1998. Around the selected Southern field only one high redshift quasar J2233--606 ($z_{em}\\simeq2.22$) has been identified (Boyle 1997),  and it is located 5.1 arcmin EAST of the HDFS. Two HST orbits have been dedicated to this object to obtain a low resolution STIS spectrum (3700 seconds with G230L and 2200 seconds with G430L). The spectrum in the range $\\lambda\\lambda=1600-5700$ \\AA~reveals, among other features, the presence of a Lyman limit break at $z\\sim1.9$, a few metal lines in the red part of the \\lya~emission line and a \\lya~forest with a few absorption lines on the top of the \\lya~emission lines. No Damped \\lya~profile is seen along the line of sight. Ground--based low resolution observations combined with the STIS spectrum have been analyzed by Sealey et al. (1998). They report a first tentative identification of metal lines and provide a HI column density of the Lyman limit system at $z=1.943$, of $N(\\rm HI) = (3.1\\pm1.0)\\times10^{17}$ \\cm2.  Here we report high spectral resolution (FWHM $\\simeq 14$ \\kms) observations obtained with the New Technology Telescope (NTT) of the nearby QSO J2233--606 ($z_{em}\\simeq2.22$) in the spectral range $\\lambda\\lambda=4386-8270$ \\AA.  Even though the 3 hours of integration time have not produced an extremely high quality spectrum, the reduced data have been released to the astronomical community (see the URL http://www.eso.org/ $\\tilde{}$ ssavagli). It is hoped that these data will provide a valuable guide when planning HDFS followup observations and will assist scientific discussions on the relation between the absorption systems along the QSO line of sight and the emitting objects which will be identified in the HDFS.  Other high resolution observations from the ground primarily covering the Lyman--$\\alpha$ forest region have been scheduled for July 1998 at the high resolution spectrograph UCLES of the Anglo Australian Telescope (AAT). ",
        "conclusions": "Unlike in the case of the HDF North program, ``follow--up'' observations of the HDFS region from the ground and space telescopes began already long before the beginning of the program. The observations from the radio to X--rays, from low and high resolution spectroscopy to narrow and broad band imaging, programmed since November 1997 to the Spring of 1999, will be combined to provide astronomers with the deepest and most complete view of the Universe for redshifts $z>1$, where most of the stars have probably already formed.  Moreover, unlike the Northern field, the HDFS is close to a high redshift quasar.  The QSO line of sight ($\\alpha = 22^h 33' 37.67''$, $\\delta = -60^\\circ 33' 28.95\"$, J2000 Equinox) is located $5'7''$ away from the Hubble Deep Field South ($\\alpha = 22^h 32' 56.22''$, $\\delta = -60^\\circ 33' 02.69\"$, J2000 Equinox). This corresponds to a real separation that is shown in Fig.~11 in the redshift range $z=0.5-3$. The separation at the QSO redshift ($z_{em}\\simeq2.22$) is 2.8 and 8.9 $h_{50}^{-1}$ Mpc ($h_{50}$ is the Hubble constant expressed in units of 50 \\kms~Mpc$^{-1}$) for a flat Universe in the physical and comoving space respectively.  Deep imaging of the region near the QSO line of sight is fundamental to the understanding the quasar environment, the origin of high redshift quasars in connection to that of galaxy clusters and the relation between \\lya~absorbers and high redshift galaxies. The typical wavelength of density fluctuations whose collapse generate galaxy clusters is around 20 $h_{50}^{-1}$ Mpc. On the other hand, once virialized, clusters have a typical size (i.e., virial radius) of about $2-5~h_{50}^{-1}$ Mpc, depending on their mass and on the background cosmological model (e.g., Kitayama \\& Suto 1996, and references therein).  At redshift $z\\sim 2.2$ we generally expect to find protoclusters which are not yet virialized. Instead, they should correspond to detectable overdensities involving scales of about 10 $h_{50}^{-1}$ Mpc comoving. Therefore, that would roughly correspond to the comoving separation between the HDFS and the main absorption clusters seen in the QSO spectrum, whereas at the same redshift, the comoving HDFS size is around 4.5 $h_{50}^{-1}$ Mpc. The size of the HFDS and the separation from the quasar might be suitable for the identification of clustered structures associated with the QSO absorption systems.  The idea of the existence of clustered structures at high redshifts has recently been confirmed by observations. Steidel et al. (1998) found evidence for a protocluster of galaxies at $z\\simeq 3.1$ in the field of two high redshift QSOs, one of which is at the same redshift as the structure. The discovered 15 Lyman break galaxies are distributed on a plane which is at least $20\\times15~h_{50}^{-2}$ Mpc$^2$ comoving. Numerical simulations have shown that non--linear structures can be present on large scales already by $z=3$. Using N--body simulations and semi--analytical methods, Governato et al.~(1998) were able to produce in the generic Cold Dark Matter scenario, structures similar to the observed one.  These high redshift associations of Lyman break galaxies are strongly biased with respect to the already overdense local dark matter distribution and are the seeds of present day galaxy clusters.  Other protoclusters of galaxies candidates at very high redshift ($z\\sim2$) have been found in a limited number of cases (Dressler et al. 1993, Francis et al. 1996, Pascarelle 1996, Hutchings 1995). At low redshift ($z<1$) extensive imaging and spectroscopic observations of galaxy clusters in quasar environments and other AGNs indicated that quasar activity is strongly correlated with the properties of the environment of the host galaxy (Yee \\& Green, 1987; Yee\\& Ellingson, 1993). The study of the HDFS QSO should allow to extend this concept to much higher redshifts. In particular, it should reveal if the highest density peaks that give origin to quasars at high redshift are also the seeds of structure formation on much larger scales."
    },
    "9804/hep-ph9804230_arXiv.txt": {
        "abstract": "For neutrinos with a magnetic moment, we show that the collisions in a  hot and dense plasma  act as an efficient  mechanism  for  the conversion of $\\nu_L$ into $\\nu_R$. The production rate for right-handed neutrinos is computed in terms of a resummed photon propagator which consistently incorporates the background effects. Assuming that the entire energy in a supernova collapse is not carried away by the $\\nu_R$, our results can be used to place an upper limit on the neutrino magnetic moment $\\mu_\\nu < (0.1-0.4)\\times 10^{-11}\\mu_B$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804042_arXiv.txt": {
        "abstract": "s{ The measurements of CMB anisotropy have opened up a window for probing the global topology of the universe on length scales comparable to and beyond the Hubble radius. We have developed a new method for calculating the CMB anisotropy in models with nontrivial topology and apply it to open universe models with compact spatial topology. We conduct a Bayesian probability analysis for a selection of models which confronts the theoretical pixel-pixel temperature correlation function with the {\\sc cobe--dmr} data. Our results demonstrate that strong constraints on compactness arise: if the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the observations.}  \\noindent The remarkable degree of isotropy of the cosmic microwave background (CMB) points to homogeneous and isotropic Friedmann-Robertson-Walker (FRW) models for the universe. This argument is a purely local one and does not refer to the global topological structure of the universe. In fact, in the absence of spatially inhomogeneous perturbations, a FRW model predicts an isotropic CMB regardless of the global topological structure. However, the observed large scale structure in the universe and CMB anisotropy allude to the existence of small spatially inhomogeneous primordial perturbations. The global topology of the universe does affect the observable properties of the CMB anisotropy.  In compact universe models, the finite spatial size usually implies a suppression of the power in large scale perturbations and consequently the CMB anisotropy is suppressed on angular scales above a characteristic angle related to size of the universe. Another signature is the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the compact manifold.  Much recent astrophysical data suggest the cosmological density parameter, $\\Omega_0$, is subcritical.~\\cite{opencase} In the absence of a cosmological constant, this would imply a hyperbolic spatial geometry for the universe (commonly referred to as the `open' universe in cosmological literature). The topologically trivial (simply connected) hyperbolic 3-space, $\\hm$, is non-compact and has infinite size.  There are numerous theoretical motivations, however, to favor a spatially compact universe.~\\cite{cct_motive} To reconcile this with a flat or hyperbolic geometry, consideration of models with non-trivial topology is required. A compact cosmological model is constructed by identifying points on the standard infinite flat or hyperbolic FRW space under the action of a suitable discrete subgroup, $\\Gamma$, of the full isometry group, $G$, of the FRW space. The FRW spatial hypersurface is the {\\em universal cover}, tiled by copies of the compact space, ${\\cal M}$. Any point ${{\\bf x}}$ of the compact space has an image ${{\\bf x}}_i = \\gamma_i {{\\bf x}}$ in each tile on the universal cover, where $\\gamma_i \\in \\Gamma$.  \\begin{figure}[b] \\plottwo{vplot_sim2.eps}{m_all_plot_sim2.eps} \\caption{The figure consists of two columns of CMB sky-maps showing a pair of $140^\\circ$ diameter hemispherical caps each, centered on the South (SGP) and North (NGP) Galactic Poles, respectively. The map labeled DATA, shows the {\\sc cobe--dmr} 53+90+31 GHz A+B data after Wiener filtering assuming a standard CDM model, normalized to {\\sc cobe}.  The rest of the five maps are one random realization of the CMB anisotropy in two examples of compact hyperbolic (CH) spaces for several values of $\\Omega_0$ based on our theoretical calculations of $C(\\hat q,\\hat q^\\prime)$ convolved with the {\\sc cobe--dmr} beam. Both surface and integrated (ISW) Sachs-Wolfe effects have been included in $C(\\hat q,\\hat q^\\prime)$. No noise was added. The power was normalized to best match the {\\sc cobe} data. The theoretical sky was optimally filtered using the {\\sc cobe} experimental noise to facilitate comparison with data.  The maps labeled L(arge)CH refer to the CH model $v3543(2,3)$. The right column shows the CMB maps for the S(mall)CH model $m004(-5,1)$. (The model number associated with the topology corresponds to that of the census of CH spaces in the Geometry center, Univ. of Minnesota; SCH is one of the smallest and LCH is one of the largest spaces in the census). LCH with $\\Omega_0=0.8$ is compatible with the data with a suitable choice of orientation while all the others are ruled out (See Table~\\ref{table1}).  For all six maps, the average, dipole and quadrupole of the $\\vert b \\vert > 20^\\circ$ sky were also removed and a $20^\\circ$ Galactic latitude cut was used, with extra cuts to remove known regions of Galactic emission proposed by the {\\sc cobe} team accounting for the ragged edges. The contours are linearly spaced at $15~\\mu {\\rm K}$ steps.  The maps have been smoothed by a $1.66^\\circ$ Gaussian filter.\\hfill\\mbox{}} \\label{fig1} \\end{figure}   For Gaussian perturbations, the angular correlation function, $C(\\hat q,\\hat q^\\prime)$, of the CMB temperature fluctuations in two directions $\\hat q$ and $\\hat q^\\prime$ in the sky completely encodes the CMB anisotropy predictions of a model.  The dominant contribution to the anisotropy in the CMB temperature measured with wide-angle beam ($\\theta_{\\sc fwhm} \\gta 2^\\circ~ \\Omega_0^{1/2}$) comes from the cosmological metric perturbations through the Sachs-Wolfe effect. The angular correlation function of the CMB anisotropy, $C(\\hat q,\\hat q^\\prime)$, depends on the spatial two point correlation function, $\\xi_\\Phi \\equiv \\langle\\Phi({\\bf x},\\taurec)\\Phi({\\bf x^\\prime},\\taurec)\\rangle $ of the gravitational potential, $\\Phi$, on the hypersurface of last scattering.\\footnote{ Other effects which contribute to the CMB anisotropy at smaller angular scales can also be approximated in terms of spatial correlation of quantities defined on the hypersurface of last scattering.~\\cite{us}} To calculate the spatial correlation function on a compact hyperbolic (CH) manifold, described by the corresponding $\\Gamma$, we have developed a general technique -- the {\\em method of images}, which evades the difficult problem \\footnote{The correlation function is usually computed using a mode function expansion. However, obtaining closed form expressions for eigenfunctions of the Laplacian may not be possible beyond the simplest topologies and even numerical estimation is known to be difficult in CH spaces.}  of solving for eigenfunctions of the Laplacian on these manifolds.~\\cite{us_texas} Using the method of images, the spatial correlation function, $\\xi^c_\\Phi$, between two points ${\\bf x}$ and ${\\bf x^\\prime}$ on a compact space of volume $V_{\\cal M}$ can be expressed as~\\cite{us} \\begin{equation} \\xic = \\lim_{r_*\\to\\infty} \\sum_{r_j < r_*} \\xiurj - \\frac{4\\pi}{V_{\\cal M}} \\int_0^{r_*} dr ~\\sinh^2r ~\\xiur\\,,~~~~~ r_j = d({\\bf x}, \\gamma_j {\\bf x}^\\prime)\\, . \\lbl{moi_final} \\end{equation} a regularized sum over the correlation function, $\\xi_\\Phi^u$, on the universal cover evaluated between ${\\bf x}$ and images $\\gamma_i{\\bf x^\\prime}$ of ${\\bf x^\\prime}$.  Numerically it suffices to evaluate the above expression up to $r_*$ a few times the curvature radius, $d_c$, to reach a convergent result. We then integrate $\\xi^c_\\Phi$ along photon trajectories to get $C(\\hat q,\\hat q^\\prime)$ which includes both surface and integrated (ISW) Sachs-Wolfe effects.  The CMB photons can be viewed as propagating to the observer from a $2$-sphere of radius, $R_H$, -- the sphere of last scattering (SLS). In contrast to the topologically trivial models, widely separated pixels in compact spaces can still have strong correlations if in the sum over images, eq.~(\\ref{moi_final}), one of the images of ${\\bf x^\\prime}$ happens to be close to ${\\bf x}$. The dependence of the spatial correlations on the anisotropic distribution of images leads to a characteristic statistical anisotropy in the CMB in compact models. These effects are pronounced when the compact space fits well within the SLS, but persist at an observable level even when $V_{\\cal M}$ is comparable to (or somewhat bigger than) $V_{\\sc sls}$, the volume of SLS.  If SLS does not fit completely inside a single tile -- a copy of the compact space, the CMB temperature values will be identical along pairs of circles if temperature fluctuations are dominated by the surface terms at the SLS.~\\cite{circles} This pattern of matched circles is one specific manifestation of the angular patterns in $C(\\hat q,\\hat q^\\prime)$.   Figure~\\ref{fig1} compares theoretical realizations of the CMB anisotropy in a selection of CH models with the {\\sc cobe--dmr} data. In this work, the six {\\sc cobe--dmr} four-year maps~\\cite{dmr4} are first compressed into a (A+B)(31+53+90 GHz) weighted-sum map, with the customized Galactic cut advocated by the {\\sc dmr} team. There is no effective loss of information when we do further data compression by using $5.2^\\circ\\times5.2^\\circ$ pixels.~\\cite{bdmr294} The theory and data maps have been postprocessed so as to facilitate a fair visual comparison. The incompatibility of models with small $V_{\\cal M}/V_{\\sc sls}$ (SCH-$\\Omega_0=0.3,0.6$) is visually obvious: the best fit amplitudes are high which is reflected in the steeper hot and cold features. Although, SCH-$\\Omega_0=0.9$ and LCH-$\\Omega_0=0.6$ do not appear grossly inconsistent, the intrinsic anisotropic correlation pattern is at odds with the data.  We have carried out a full Bayesian analysis of the probability of the CH models given the {\\sc cobe--dmr} 4yr data. In Table~\\ref{table1} we present the {\\em relative likelihood} of the selected models to that of the infinite, $\\hm$, model with the same $\\Omega_0$. (The {\\sc cobe} data alone does not strongly differentiate between the infinite hyperbolic models with different $\\Omega_0$.) The anisotropy of the theoretical $C(\\hat q,\\hat q^\\prime)$ causes the likelihood of compact models to vary significantly with the orientation of the space with respect to the sky, depending on how closely the features in the single data realization available match (or mismatch) the pattern in $C(\\hat q,\\hat q^\\prime)$. Some optimal orientations may also have the ``ugly'' correlation features hidden in the Galactic cut.  We analyzed $24$ different orientations for each of our models and found that only the model with $V_{\\cal M} > V_{\\sc sls}$ (LCH-$\\Omega_0=0.8$) cannot be excluded (at one orientation this model is even preferable to standard CDM; this raises a question of the statistical significance of any detection of intrinsic anisotropy of a space when only a single realization of data is available).  Similar conclusions were reached by some of the authors (JRB, DP and I. Sokolov~\\cite{us_torus}) for flat toroidal models.  Comparison of the full angular correlation with {\\sc cobe} data led to a much stronger limit on the compactness of the universe than limits from other methods.~\\cite{tor_refs} The main result of the analysis was that $V_{\\sc sls}/V_{\\!\\cal M} < 0.4$ at $95\\%~CL$ for the equal-sided $3$-torus. For non compact $1$-torus, the constraint on the most compact dimension is not quite as strong.  In summary, our results demonstrate that the {\\sc cobe} data can put strong constraints on the compact models of the universe. If the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the large angle {\\sc cobe--dmr} data. \\vspace*{-2mm} \\begin{table}[htb] \\caption{The Log-likelihoods of the compact hyperbolic models relative to the infinite models with same $\\Omega_0$ are listed below.  The likelihoods are calculated by comparison with {\\sc cobe--dmr} data. The three columns of Log-likelihood ratios correspond to the best, second best and worst values that we have obtained amongst $24$ different rotations of the compact space relative to the sky. The number in brackets gives a convenient, albeit crude, translation to gaussian likelihood. Only the last model can be reconciled with the {\\sc cobe--dmr} data.\\hfill\\mbox{}\\label{logprob_tab}} \\vspace{0.3cm} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline Topology &$\\Omega_0$&$V_{\\sc sls}/V_{\\!\\cal M}$&\\mco{3}{|c|}{Relative Log. Likelihood (Gaussian approx.)}\\\\ \\cline{4-6} & & &\\mco{3}{|c|}{Orientation}\\\\ &    &    &`best'&`second best'&`worst'\\\\ \\hline $m004(-5,1)$ &0.3&153.4 & -35.5 (8.4$\\sigma$)& -35.7 (8.4$\\sigma$) & -57.9 (10.8$\\sigma$)\\\\ &0.6&19.3 & -22.9 (6.8$\\sigma$)& -23.3 (6.8$\\sigma$) & -49.4 ( 9.9$\\sigma$)\\\\ $V_{\\cal M}/d_c^3=0.98$      &0.9&1.2  & -4.4  (3.0$\\sigma$)& -8.5  (4.1$\\sigma$) & -37.4 ( 8.6$\\sigma$)\\\\ \\hline $v3543(2,3)$ & 0.6 &2.9 & -3.6 (2.7$\\sigma$) & -5.6 (3.3$\\sigma$)& -31.0 (7.9$\\sigma$)  \\\\ $V_{\\cal M}/d_c^3=6.45$ & 0.8 &0.6 &  2.5 (2.2$\\sigma$) & -0.8 (1.3$\\sigma$)   & -12.6 (5.0$\\sigma$)  \\\\ \\hline \\end{tabular} \\end{center} \\lbl{table1} \\end{table}  \\vspace*{-6mm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804274_arXiv.txt": {
        "abstract": "We search for angular correlation of gamma-ray bursts with cataloged quasars, BL Lac objects, and AGN using a large sample of relatively well-localized bursts detected by WATCH on board GRANAT and EURECA, IPN, and BATSE (327 bursts total). A statistically significant (99.99\\% confidence) correlation between GRB and $M_B<-21$ AGN in the redshift range $0.1<z<0.32$ is found.  The correlation with AGN is detected, with a lower significance, in three independent GRB datasets. The correlation amplitude implies that, depending on the AGN catalog completeness, 10\\% to 100\\% of bursts with peak fluxes in the range $3-30\\times10^{-6}$~erg~s$^{-1}$~cm$^{-2}$ in the 100--500 keV band are physically related to AGN. The established distance scale corresponds to the energy release of order $10^{52}$~ergs per burst. ",
        "introduction": "Gamma-ray bursts are distributed isotropically on the sky; their peak flux distribution shows the lack of faint bursts compared to the expectation for a homogeneous distribution in Euclidean space (see e.g.\\ Fishman \\& Meegan 1995 for review). These observed properties of GRB are reproduced by two popular models: 1) GRB arise in an extended Galactic halo with a core radius of $\\sim 100$~kpc and 2) GRB are located at cosmological distances.  The Galactic halo models are challenged by the observed isotropy of the burst positions. The isotropy constraints much improved recently, and as a result, most variants of the halo model are no longer viable (Briggs et al.\\ 1996).  An additional argument against the Galactic halo model comes from {\\em Einstein} data. The {\\em Einstein} IPC sensitivity is sufficient to detect bursts from the halos of nearby galaxies, but they were not found in the data (Hamilton et al.\\ 1996).  In the cosmological model, the burst isotropy and departures of $\\log N - \\log S$ from the Euclidean $S^{-3/2}$ law are explained naturally.  The minimum redshift $z=0.835$ of the optical transient associated with GRB~970528 (Metzger \\etal\\ 1997) is a decisive evidence in favor of the cosmological model.  Optical spectroscopy of the GRB counterparts is complicated by the faintness of associated optical transients and their relatively featureless spectrum (e.g., no strong emission lines were observed in the GRB~970528 optical transient).  Therefore, indirect estimates of the GRB distance scale are still useful.  The shape of the GRB peak flux distribution implies that sources of the dimmest BATSE bursts are at $z\\sim1$, assuming no source evolution (Emslie \\& Horack 1994; Fishman \\& Meegan 1995 and references therein). For some plausible evolution of GRB volume density, the above distance estimate can vary by a factor of $\\gtrsim 2$ (Horack et al.\\ 1995). Using a different approach, Quashnock (1996) derived $z>0.25$ for the dimmest BATSE bursts by cross-correlating the third BATSE catalog with the known large scale structure at low redshifts.  Some earlier studies already searched for a direct relationship between GRB and other astrophysical objects at cosmological distances. Kolatt \\& Piran (1996) and Marani et al.\\ (1997) found that GRB from the third BATSE catalog are correlated with Abell clusters. However, using more precise GRB localizations, Burenin et al.\\ (1997), Hurley et al.\\ (1997), and Gorosabel \\& Castro-Tirado (1997) have not found any correlation with Abell clusters. Furthermore, the amplitude of correlation of well-localized GRB and $z<0.1$ Abell clusters is lower than expected in the case of the same spatial distribution of GRB and optically luminous matter (Burenin et al.\\ 1997). If this result is interpreted as an indication that bright and well-localized GRB are at greater distances than nearby Abell clusters, the dimmest BATSE bursts should be at $z>0.3$, in agreement with Quashnock's (1996) results.  At still higher redshifts, $z\\sim0.1-1$, a natural choice is to search for correlation of GRB with quasars and other flavors of active galactic nuclei, which comprise the majority of known objects at these high redshifts. To search for correlation of AGN and GRB is also attractive because some theoretical models relate bursts to physical processes in active nuclei (e.g.\\ Lejter 1980, Carter 1992). Some attempts to search for such a correlation have been undertaken earlier. A marginal evidence for excess of QSO in the small Interplanetary Network (IPN) GRB error boxes has been found (Vrba \\etal\\ 1995). However, Webber \\etal\\ (1995) and Gorosabel \\etal\\ (1995) found that the number of QSO and AGN in GRB error boxes is consistent with random. These earlier analyses were based on a smaller number of bursts with good localizations than available at present.  Citing a recent work, Schartel \\etal\\ (1997) have found a correlation of well-localized GRB from the third BATSE catalog with radio-quiet QSO. The strongest correlation was detected at the $>99.7\\%$ confidence level for intrinsically bright QSO with $z<1$.  The goal of this work is to search for a correlation of GRB with QSO and AGN using all available data for good, with $\\lesssim 1^\\circ$ uncertainty, bursts localizations. We use $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0=0$.  \\begin{figure*}[htb] \\begin{minipage}[t]{0.47\\textwidth} \\epsfysize=1.1\\textwidth \\centerline{\\epsffile{fig1.ps}} \\vskip -5pt \\small {\\bf Fig.~1}--- Number density of QSO from VCV96 as a function of redshift. Smooth solid line represents the true number density of these objects (accounting for their cosmological evolution, Franceschini \\etal\\ 1994). \\end{minipage} \\hfill \\begin{minipage}[t]{0.47\\textwidth} \\epsfysize=1.1\\textwidth \\centerline{\\epsffile{fig2.ps}} \\vskip -5pt \\small {\\bf Fig.~2}--- Number density of AGN from VCV96 as a function of their redshift.  Dotted histogram represents number density of all AGN, the solid one corresponds to bright $M_B<-21$ AGN. \\end{minipage} \\end{figure*} ",
        "conclusions": "This is not the first attempt to find optical counterparts of GRB by analyzing the optical content of small-area localizations, so the comparison of our results with some previous works is in order.  Vrba et al.\\ (1995) have performed extensive optical photometry of 8 small-area ($<70$~arcmin$^2$) IPN localizations searching for objects with unusual colors, variability, and proper motions. Only blue objects, which Vrba et al.\\ interpreted as QSO, showed a marginal excess at a rate approximately one per localization. This is in agreement with our results.  Webber et al.\\ (1995) and Gorosabel et al.\\ (1995) used several tens GRB localizations from IPN and WATCH, respectively. The optical content was taken from existing catalogs, similar to our study. No excess of either AGN or any other class of objects was found. These results do not directly contradict to ours because Webber et al.\\ (1995) and Gorosabel et al.\\ (1995) used a smaller number of bursts, did not account for the incompleteness of optical catalogs (\\S2.2), and did not apply redshift constraints in search for correlation. Since the completeness of the VCV96 catalog (and probably other all-sky catalogs) is below, or on the level of, several percent (\\S4), analyses based on several tens of bursts are inconclusive.   The most similar to ours is the work of Schartel et al.\\ (1997) who cross-correlated BATSE bursts with the AGN and QSO from VCV96. They found a marginal correlation with QSO, and no correlation with AGN, with no redshift constraints applied to AGN. We essentially reproduce these results (Fig.~4 and Table~1). A significant correlation with AGN found in our analysis is not found by Schartel et al.\\ because they used a smaller number of bursts, with poorer localizations (only BATSE data), and used the entire VCV96 catalog which is very incomplete for AGN even at low redshift (\\S2.2).  Finally, we mention the two GRB with small-area localizations, in which AGN were found. Drinkwater et al.\\ (1997) report that the possible X-ray counterpart of GRB~920501 is associated with a Seyfert~1 galaxy at $z=0.315$. The first X-ray localization of a gamma-ray burst by BeppoSAX (3$^\\prime$ radius) contains a $z=1.038$ QSO (Piro et al.\\ 1998). On the other hand, optical transients associated with other BeppoSAX bursts, GRB~970508, 970228, and 971214 probably are not AGN.  Our analysis supersedes most of earlier searches for GRB-AGN association because we used a large GRB dataset, carefully accounted for incompleteness of the optical catalogs, and introduced sensible object redshift constraints. These advantages made it possible to find the strongest ever evidence for association of GRB with a known class of extragalactic objects. However, there are several problems with our analysis. First, we used relatively large area gamma-ray burst localizations, and therefore had to perform a statistical correlation rather than an object-by-object identification.  Second, we used a sparse catalog of AGN and QSO, in which many objects were found in a number of small-area high-sensitivity searches. In fact, this introduces a possibility that the correlation we detect is not with AGN but rather with objects around which the AGN we searched, e.g.\\ normal galaxies or other QSO (data from Arp 1980 and Monk et al.\\ 1988), targets of {\\em Einstein}\\/ pointings (EMSS AGN, Stocke et al.\\ 1991). However, the sky distribution of GRB with $0.1<z<0.32$ AGN inside the localization area, does not generally follow the regions of deep AGN surveys (Fig.~7). Also, a clear redshift dependence of the correlation is not easily explainable in such a scenario.  Another problem is that the amplitude of the detected correlation implies that the fraction of bursts related to AGN is somewhat higher than the estimated completeness of the optical catalog (\\S4). We can offer two possible explanations. First is that GRB prefer luminous AGN, for which the catalog completeness is higher. Second is that GRB prefer a certain type of AGN (Seyfert 1 or 2, X-ray loud or quiet, radio loud or quiet, etc.) which is more commonly present in the VCV96 catalog than the ``average'' AGN.  Despite these problems, we believe that the case for an association of bright gamma-ray bursts with AGN at moderate redshift is compelling.  This case can be further proved, or disproved, by an extensive optical, preferably spectroscopic, survey of small-area localizations of bright GRB, similar to the work of Vrba et al.\\ (1995) but using a larger number of bursts. Another approach would be to use a more complete catalog of AGN covering a significant fraction of the sky; unfortunately, this seems impractical until the Sloan Digitized Sky Survey is completed."
    },
    "9804/astro-ph9804104_arXiv.txt": {
        "abstract": "Recently, \\citet{MParkJGott97a} claimed that there is a statistically significant, strong, negative correlation between the image separation $\\Delta\\theta$ and source redshift $\\zs$ for gravitational lenses.  This is somewhat puzzling if one believes in a flat ($k=0$) universe, since in this case the typical image separation is expected to be independent of the source redshift, while one expects a negative correlation in a $k=-1$ universe and a positive one in a $k=+1$ universe. \\citeauthor{MParkJGott97a} explored several effects which could cause the observed correlation, but no combination of these can explain the observations with a realistic scenario.  Here, I explore this test further in three ways.  First, I show that in an inhomogeneous universe a negative correlation is expected regardless of the value of $k$. Second, I test whether the \\thz\\ can be used as a test to determine $\\lnull$ and $\\onull$, rather than just the sign of $k$.  Third, I compare the results of the test from the \\citeauthor{MParkJGott97a} sample to those using other samples of gravitational lenses, which can illuminate (unknown) selection effects and probe the usefulness of the \\thz\\ as a cosmological test. ",
        "introduction": "Historically, there has been little interest in the \\thz\\ compared to other cosmological tests based on gravitational lensing statistics, perhaps because the inflationary paradigm \\citep[\\eg][]{AGuth81a}, which began about the same time as the discovery of the first gravitational lens \\citep*{DWalshCW79a}, has become so influential.  Since a flat ($k=0$) universe is a robust prediction of inflation, many researchers assume this and consider only flat universes (or, at most, $k=-1$ cosmological models with $\\lnull=0$).  Due to the fact that for the popular singular isothermal sphere model for a single-galaxy lens the average image separation $\\Delta\\theta$, integrated over the lens redshift $\\zd$ from $\\zd=0$ to $\\zd=\\zs$, is {\\em completely independent\\/} of the source redshift $\\zs$ in a flat universe, there is little point in pursuing the \\thz\\ if one is interested primarily in flat cosmological models.  If one is not committed to a flat universe, then of course one should not assume $k=0$, but even if one believes that the universe must be flat, it is still important to test this belief observationally.  The situation is somewhat worsened by the fact that most `standard' cosmological tests such as the \\mbox{$m$-$z$} (magnitude-redshift or `standard candle') and $\\theta$-$z$ (angular size-redshift or `standard rod') relations, `conventional' gravitational lensing statistics, age of the universe) are relatively insensitive to the radius of curvature of the universe ($R_{0} \\sim (|\\Omega_{0}+\\lambda_{0}-1|)^{-\\frac{1}{2}}$), being degenerate in combinations of $\\lnull$ and $\\onull$ in directions roughly perpendicular to lines of constant $R_{0}$ in the $\\lnull$-$\\onull$ plane.  A notable exception are constraints derived from CMB anisotropies \\citep*[\\eg][]{DScottSW95a,WHuSS97a}. ",
        "conclusions": "\\citet{MParkJGott97a} pointed out that the image separations in gravitational lens systems show a strong significant negative correlation with the source redshift, while in a flat universe one would expect no correlation (while a negative correlation would be expected in a universe with negative curvature and a positive one in a universe of positive curvature).  None of the possibilities they examined were strong enough to explain the effect.  A possibility not examined by them, namely an inhomogeneous universe, produces a negative correlation regardless of the sign of the curvature, but it too is not strong enough to account for the effect.  As a general test for the values of $\\lnull$ and $\\onull$ the test is of no use, all cosmological models being assigned roughly the same probability, but {\\em which\\/} value they are assigned depends on the sample used.  The strong dependence of the result on the sample used seems to indicate that the result of \\citet{MParkJGott97a} is due not to some physical cause but rather to unidentified selection effects in the sample of gravitational lenses taken from the literature.  The large number of JVAS and CLASS lenses gives us an independent comparison sample, thus demonstrating the need for discovering a large number of lenses in a well-defined sample.  As \\citet{MParkJGott97a} point out, since many conclusions based on `conventional' gravitational lensing statistics are based on essentially the same lenses as in their literature sample, if this sample is for some unknown reason atypical, then conclusions drawn from statistical analyses of it must be examined with care.  It will thus be interesting to see what conclusions can be drawn from a statistical analysis of the JVAS/CLASS sample after the observational tasks have been completed.  (We expect to find more lenses, but have no qualms about using the present incomplete sample in this analysis since there is no reason to believe that a larger sample would show a different \\thz.)"
    },
    "9804/astro-ph9804038_arXiv.txt": {
        "abstract": "We use numerical simulations of structure formation in a Cold Dark Matter model to predict the absorption lines in the soft X-rays produced by heavy elements in the shock-heated intergalactic medium at low redshift. The simulation incorporates a model for heavy element production in galaxies and the subsequent dispersion of the metals to the intergalactic medium. We analyze in particular absorption lines produced by oxygen, and calculate the ionization stage taking into account the observed X-ray background at the present time. We find that oxygen is fully ionized by the X-ray background in low-density voids, and is mostly in the form of $\\OVII$ and $\\OVIII$ in the sheets and filamentary regions. Strong absorption lines of $\\OVII$ and $\\OVIII$ with equivalent widths $W\\sim 100 \\kms$ are produced in filamentary regions of overdensities $\\sim 100$ and temperatures $\\sim 10^6$ K, located in the outskirts of groups and clusters of galaxies. The $\\OVII$ line at $E=574$ eV is generally the strongest one in these systems. Our model predicts that any X-ray source (such as a quasar) should typically show about one $\\OVII$ absorption line with $W > 100 \\kms$ in the interval from $z=0$ to $z=0.3$. These lines could be detected with the upcoming generation of X-ray telescopes, and their origin in intervening systems could be confirmed by the association with groups of galaxies and X-ray emitting halos near the line-of-sight at the same redshift. The hot intergalactic medium may be one of the main reservoirs of baryons in the present universe, and the heavy element X-ray absorption lines offer a promising possibility of detecting this new component in the near future. ",
        "introduction": "Hierarchical theories of the formation of large-scale structure are based on the hypothesis that the gravitational collapse of some type of cold dark matter (i.e., any collisionless matter with a small enough velocity dispersion to allow collapse down to scales much smaller than the observed structure), triggered by initial density fluctuations, is responsible for the formation of galaxies and for the later assembly of groups, clusters, and superclusters of galaxies. This generic scenario has been very successful in explaining a large variety of observations, even though the origin of the initial density fluctuations remains unknown, and simple models are therefore assumed for the primordial power spectrum (generally based on adiabatic, Gaussian, scale-invariant fluctuations).  One of the predictions that have been made from these models, with the use of numerical hydrodynamic simulations, is that despite the wide range of scales over which the baryonic matter is able to collapse (as dark matter halos successively merge from the onset of non-linearity until the present time), a relatively large fraction of baryons should still remain as ``intergalactic matter''. This intergalactic medium should be forming a network of shock-heated gas in the form of filamentary and sheet-like structures connected to galaxy clusters and groups, as well as colder gas left out in voids, as is revealed in numerical simulations (e.g., Ostriker \\& Cen 1996). A similar network of photoionized and shock-heated gas should also be present at high redshift (although at lower temperatures than at the present time, due to the lower velocities of collapse at high redshift), and probably gives rise to the hydrogen \\lya forest (e.g., Cen \\etal 1994, Hernquist \\etal 1996, Miralda-Escud\\'e \\etal 1996, Zhang \\etal 1997, 1998).  Here, a new prediction based on the same theory of hierarchical clustering and the presence of intergalactic gas shall be studied. It has already been shown observationally that heavy elements are present in the high-redshift absorption systems (Songaila \\& Cowie 1996 and references therein). It is very likely that many more elements were spread to the intergalactic medium (hereafter, IGM) at later times, either through galactic winds energized by supernova explosions or active galactic nuclei (e.g., Dekel \\& Silk 1986), or simply by the process of gravitational merging, which can also lead to some gas being ejected from halos back into the IGM (Gnedin \\& Ostriker 1997). As pointed out by Shapiro \\& Bahcall (1980) and Aldcroft \\etal (1994), heavy elements in the intergalactic medium should cause absorption lines (as well as continuum edges of absorption) on background X-ray sources In fact, for highly ionized, hot gas at $T\\sim 10^6\\kelvin$, absorption lines from heavy elements are probably the only method of detection, since hydrogen is highly ionized and its absorption lines are very weak. The soft X-ray emission from low-density gas is also very weak and would generally be seen superposed with emission from other structures along the line-of-sight, as well as the Galactic emission. We shall denote these absorption lines by ``X-ray forest'', in analogy to the \\lya forest caused by hydrogen. The X-ray forest is much more difficult to observe than the \\lya forest, due to the lower sensitivity and resolution in the X-ray band. The possibility of observing this X-ray forest has also been discussed by Aldcroft \\etal (1994) and Canizares \\& Fang (1998).  Similar absorption lines may also be detectable in the ultraviolet when the temperature is not very high; these are generally caused by lithium-like ions like $\\OVI$. Mulchaey \\etal (1996) proposed this as a method to detect halos of hot has in poor groups of spiral galaxies, where the temperature of the halo gas may be too low to have been detected in emission by ROSAT. There is in fact evidence for a population of $\\OVI$ absorbers among the numerous absorption systems seen in quasar spectra that may arise in hot, collisionally ionized gas (e.g., Burles \\& Tytler 1996).  In this paper we shall predict some properties of the X-ray forest from a hydrodynamic simulation, and discuss the prospects for detecting it in future X-ray missions. ",
        "conclusions": "We have presented in this paper a new generic prediction of the large-scale structure theories based on hierarchical gravitational collapse of density fluctuations: X-ray quasars should show absorption lines in their spectra from heavy elements in the intervening intergalactic gas. A substantial fraction of the baryons are predicted to be in the form of low-density intergalactic gas at the present time (Ostriker \\& Cen 1996; Cen \\& Ostriker 1998; Miralda-Escud\\'e \\etal 1996; Zhang \\etal 1998). Several mechanisms are known to enrich this gas with heavy elements; given the observations of the metallicity in the centers of rich clusters, it seems inevitable that this enrichment has also taken place in the lower-density gas. In fact, the simulations presented here {\\it underpredict} the metal abundance observed in clusters; had we increased the metal injection rate to fit the present cluster metallicities, the column densities of the absorbers we predict would be proportionally increased.  Absorption lines in the soft X-rays, as well as far-UV lines produced by lithium-like ions (Mulchaey \\etal 1996), are the only observational means we know of to detect low-density, hot intergalactic gas in the present universe. At temperatures $\\sim 10^6$ K, the neutral hydrogen fraction is too low to produce significant \\lya absorption, so we must rely upon the highly ionized heavy elements. The search for intergalactic gas is important to complete an inventory of the baryons in the present universe. The total baryon density observed in galaxies and X-ray emitting gas in clusters comes close to the baryon density predicted by nucleosynthesis (e.g., Persic \\& Salucci 1997; Fukugita, Hogan, \\& Peebles 1998), but given the uncertainties it is possible for up to $\\sim$ 80\\% of the baryons to be undetected, in the form of ionized intergalactic gas. Results from cosmological simulations also suggest that the amount of baryons in the \\lya forest at $z\\sim 3$ is higher than the sum of all the baryons observed at present (Rauch \\etal 1997; Weinberg \\etal 1998; Zhang \\etal 1998).  The detection and subsequent study of the X-ray forest would lead to a large number of potential applications. The number of absorption lines found will be measuring the product of the baryon density times the metallicity in the IGM. The ratio of strengths of the $\\OVII$ and $\\OVIII$ lines will probe the distribution of the gas temperature and density, and several other lines may be observed which could provide extensive tests for the predictions of large-scale structure models. The X-ray absorption lines could be correlated with structures seen in emission near the line-of-sight, such as galaxies or X-ray emitting clusters and groups. A new era in the study of the intergalactic medium could be opened with the discovery of the X-ray forest by the new X-ray observatories."
    },
    "9804/hep-ph9804205_arXiv.txt": {
        "abstract": "s{One of the fundamental problems of modern cosmology is to explain the origin of all the matter and radiation in the Universe today. The inflationary model predicts that the oscillations of the scalar field at the end of inflation will convert the coherent energy density of the inflaton into a large number of particles, responsible for the present entropy of the Universe. The transition from the inflationary era to the radiation era was originally called reheating, and we now understand that it may consist of three different stages: preheating, in which the homogeneous inflaton field decays coherently into bosonic waves (scalars and/or vectors) with large occupation numbers; backreaction and rescattering, in which different energy bands get mixed; and finally decoherence and thermalization, in which those waves break up into particles that thermalize and acquire a black body spectrum at a certain temperature. These three stages are non-perturbative, non-linear and out of equilibrium, and we are just beginning to understand them. In this talk I will concentrate on the preheating part, putting emphasis on the differences between preheating in chaotic and in hybrid inflation. } ",
        "introduction": "At the end of inflation all the energy density is in the homogeneous zero mode of the inflaton field. The Universe is in a vacuum-like state with zero temperature and vanishing particle and entropy densities. The problem of reheating is how to convert all this coherent energy into a state of thermalized relativistic particles. The original analysis~\\cite{book} assumed the perturbative decay of the inflaton into bosons and fermions, as if the inflaton were already an ensemble of decoherent {\\em particles}. Reheating ended when the total decay rate of the inflaton was of the order of the expansion rate of the Universe, $\\Gamma\\sim H$, while the total energy of the inflaton field decayed exponentially fast into other particles. As a consequence, the final reheating temperature only depended upon $\\Gamma$. We understand today that, for certain parameter ranges, there is a new decay channel that is non-perturbative,\\cite{KLS1} due to the coherent oscillations of the inflaton field, which induces stimulated emission of bosonic~\\footnote{Fermions can also be parametrically amplified, but their occupation numbers are constrained, $n_k\\leq1$, by Pauli's exclusion principle.} particles into energy bands with large occupation numbers. The modes in these bands can be understood as Bose condensates, and they behave like classical waves. The backreaction of these modes on the homogeneous inflaton field and the rescattering among themselves produce a state that is far from thermal equilibrium and may induce very interesting phenomena, such as non-thermal phase transitions~\\cite{nonthermal} with production of topological defects, a stochastic background of gravitational waves,\\cite{GW} production of heavy particles in a state far from equilibrium, which may help GUT baryogenesis~\\cite{baryo} or constitute today the dark matter in our Universe.\\cite{DM} These classical waves eventually reach a state of turbulence where, hopefully, decoherence will occur and thermalization will follow, although these stages are not yet fully understood, either analytically or numerically.  The period in which particles are produced via parametric resonance is called {\\em preheating}.\\cite{KLS1} The idea is relatively simple, the oscillations of the inflaton field induce mixing of positive and negative frequencies in the quantum state of the field it couples to. In the language of quantum fields in curved space, creation and annihilation operators mix via Bogoliubov transformations,\\footnote{For an introduction to particle production in strong external fields, see Grib et al.\\cite{GMM}} $a_k = \\alpha_k \\bar a_k + \\beta^\\ast \\bar a_{-k}^\\dagger$, and with every oscillation of the inflaton field, new particles are produced, $n_k \\equiv \\langle\\bar 0|a_k^\\dagger a_k|\\bar 0\\rangle = |\\beta_k|^2$. In the case of chaotic inflation, with a massive inflaton $\\phi$ coupled to a massless scalar field $\\chi$, the evolution equation for the Fourier modes, $\\ddot X_k + \\omega_k^2 X_k=0$, with $X_k=a^{3/2}(t)\\chi_k$ and $\\omega_k^2 = k^2/a^2(t) + g^2\\phi^2(t)$, can be cast in the form of a Mathieu equation, with coefficients $A=k^2/4a^2m^2+2q$ and $q=g^2\\Phi^2/4m^2$, where $\\Phi$ is the amplitude and $m$ is the frequency of inflaton oscillations, $\\phi(t)=\\Phi(t)\\sin mt$. For certain values of the parameters $(A,q)$ there are exact solutions that grow exponentially with time, and each mode $k$ belongs to an instability band of the Mathieu equation.\\footnote{For a recent comprehensive review on preheating after chaotic inflation see Kofman~\\cite{Kofman} and references therein.} These instabilities can be interpreted as coherent ``particle'' production with large occupation numbers. One way of understanding this phenomenon is to consider the energy of these modes as that of a harmonic oscillator, $E_k = |\\dot X_k|^2/2 + \\omega_k^2 |X_k|^2/2 = \\hbar \\omega_k (n_k + 1/2)$.  The occupation number of level $k$ can grow exponentially fast, $n_k\\sim\\exp(2\\mu_k mt)\\gg1$, and these modes soon behave like classical waves. It is analogous to the well-known mechanism of generation of density perturbations during inflation.\\cite{book} The parameter $q$ during preheating determines the strength of the resonance. It is possible that the model parameters are such that parametric resonance does {\\em not} occur, and then the usual perturbative approach would follow, with decay rate $\\Gamma$. In fact, as the Universe expands, the growth of the scale factor and the decrease of the amplitude of inflaton oscillations shifts the values of $(A,q)$ along the stability/instability chart of the Mathieu equation, going from broad resonance, for $q\\gg1$, to narrow resonance, $q\\ll1$, and finally to the perturbative decay of the inflaton.  Parametric resonance will stop whenever the inflaton decay is dominated by the perturbative decay, $q m < \\Gamma$, or when the instability modes are redshifted away from the (last) narrow resonance band, $q^2 m < H$. ",
        "conclusions": ""
    },
    "9804/hep-ph9804425_arXiv.txt": {
        "abstract": "During preheating after inflation, parametric resonance rapidly generates very large fluctuations of scalar fields. In models where the inflaton field $\\phi$ oscillates in a double-well potential and interacts with another scalar field $X$, fluctuations of $X$ can keep the $\\phi\\to-\\phi$ symmetry temporarily restored. If the coupling of $\\phi$ to $X$ is much stronger than the inflaton self-coupling, the subsequent symmetry breaking is a first-order phase transition. We demonstrate the existence of  this nonthermal  phase transition with lattice simulations of the full nonlinear dynamics of the interacting fields. In particular, we observe nucleation of an expanding bubble. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804326_arXiv.txt": {
        "abstract": "{ We study the galaxy number density (smoothed on a $5h^{-1}$~Mpc scale) around 18 low-redshift Lyman-alpha absorbers previously observed with {\\sl HST}.  The absorbers lie in the foregrounds of Mrk 335, Mrk 421, Mrk 501, I Zw 1, and 3C 273, all within regions where there are now complete redshift surveys to $m_{\\rm Zw}=15.5$.  We construct a smoothed galaxy number density field from the redshift survey data and determine the distribution of densities at the Lyman-alpha absorber locations.  We also find the distribution of galaxy number density for a variety of test samples: all galaxy locations within the Center for Astrophysics Redshift Survey (CfA2), CfA2 galaxy locations along randomly selected lines of sight, and randomly chosen redshifts along random lines of sight.  The Lyman-alpha absorbers are present in dense regions of the survey, but occur far more frequently in underdense regions than do typical luminous galaxies.  The distribution of smoothed galaxy density around the Lyman-alpha absorbers is inconsistent at the 4$\\sigma$ level with the density distribution around survey galaxies.  It is highly consistent with a density distribution at randomly chosen redshifts along random lines of sight.  This supports earlier evidence that the nearby, low column density ($\\log N_{H\\sc i} \\lesssim 14$) Lyman-alpha forest systems are spatially distributed at random; they are not well correlated with the local large-scale structure. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804110_arXiv.txt": {
        "abstract": "We show that the luminosity dependence of the red clump stars on age and metallicity can cause a difference of up to $\\lesssim$0.6 mag in the mean absolute I magnitude of the red clump between different stellar populations. We show that this effect may resolve the apparent $\\approx$0.4 mag discrepancy between red clump-derived distance moduli to the Magellanic Clouds and those from, e.g., Cepheid variables. Taking into account the population effects on red clump luminosity, we determine a distance modulus to the LMC of 18.36 $\\pm$ 0.17 mag, and to the SMC of 18.82 $\\pm$ 0.20 mag.  Our alternate red clump LMC distance is consistent with the value (m-M)$_{LMC}$ = 18.50 $\\pm$ 0.10 adopted by the HST Cepheid Key Project.  We briefly examine model predictions of red clump luminosity, and find that variations in helium abundance and core mass could bring the Clouds closer by some 0.10--0.15 mag, but not by the $\\approx$0.4 mag that would result from setting the mean absolute I-magnitude of the Cloud red clumps equal to the that of the Solar neighborhood red clump. ",
        "introduction": "The distance to the Magellanic Clouds is a problem of great astrophysical interest due to their role in the determination of extragalactic distances (c.f. \\cite{mad98}). Despite the large amount of effort that has gone into the determination of these distances, they remain a matter of some controversy (see \\cite{wes97}, \\cite{wes90} for thorough discussions).  The success of the HST Cepheid Key Project, whose goal is the determination of H$_0$ to an accuracy of 10\\%, depends critically on knowledge of the distance to the Clouds, especially the LMC.  Based on the Cepheid variable Period-Luminosity relation (\\cite{mad98}) and the light echoes of SN1987A (\\cite{pan97}), the Key Project has adopted a distance modulus of (m-M)$_{LMC}$ = 18.50 $\\pm$ 0.10 (50 $\\pm$ 2 kpc) (\\cite{raw97}). These determinations are in good agreement with recent derivations from RR Lyrae variables, based on both ground-based, statistical parallax methods (\\cite{fea97b}), and on the {\\sl Hipparcos}-calibrated distance scale (\\cite{gra97}, \\cite{rei97}).  However, the LMC distance controversy is far from settled.  A substantial fraction of recent techniques have yielded smaller values, typically in the range of (m-M)$_{LMC}$ $\\approx$ 18.30 $\\pm$ 0.10.  These include further analysis of the SN1987A light echoes (\\cite{gou98}) and recalibration of the RR Lyrae magnitude-metallicity relation (\\cite{lay96}).  If the distance to the LMC is uncertain, that of the SMC is even more so (\\cite{wes97}, \\cite{uda98}).  Because of its smaller Cepheid population and its large line-of-sight depth, distance determinations to the SMC are in generally poorly constrained.  We can only say with confidence that it lies some 0.3--0.6 mag beyond the LMC.  The best Cepheid distance to the SMC is 18.94 $\\pm$ 0.04 (\\cite{lan94}).  Recently, the large photometric sample of the OGLE microlensing survey and its resulting high-quality color-magnitude diagrams of the Clouds have permitted the development of the red clump of intermediate-age helium-burning stars as a ``standard candle'' for single-step distance determinations (\\cite{pac98}). The key to this method has been the availability of accurate {\\sl Hipparcos} parallaxes to calibrate the absolute magnitude of solar neighborhood red clump stars.  The {\\sl Hipparcos} color-magnitude diagrams (e.g., \\cite{jim98}) show a solar-neighborhood red clump that has a very small dispersion in mean absolute I-band magnitude.  Because the red clump is the dominant post-main-sequence evolutionary phase for most stars, it makes a tempting target for the application of ``standard candle'' techniques for distance determinations.  The red clump method was developed very thoroughly and applied to fields in the Galactic bulge (\\cite{pac98}) and M31 (\\cite{sta98}). These studies used the mean absolute magnitude of solar-neighborhood red clump stars, M$_I^0$ = $-$0.23 $\\pm$ 0.03, obtained from a volume-limited sample of 228 red clump stars observed with {\\sl Hipparcos}. The red clump method is based on a well-populated, well-calibrated phase of stellar evolution; in fact, it is possibly a more reliable distance indicator than many other methods that have been employed.  In a recent paper, \\cite{uda98} extended the red clump method to the Magellanic Clouds.  After taking careful account of the reddening distributions along Cloud lines of sight, they find mean red clump magnitudes of I$_0$(LMC) = 17.85 $\\pm$ 0.03, and I$_0$(SMC) = 18.33 $\\pm$ 0.03.  Using the solar neighborhood value M$_I^0$ = $-$0.23 $\\pm$ 0.03, \\cite{uda98} find distance moduli of (m-M)$_{LMC}$ = 18.08 $\\pm$ 0.15 and (m-M)$_{SMC}$ = 18.56 $\\pm$ 0.09. These values are $\\approx$0.4 mag below the ``long'' distance scale preferred by the HST Cepheid Key Project, and only marginally consistent with the ``short'' scale.  \\cite{sta98b} applied the same technique to an independent large photometric survey of the LMC and obtained the virtually identical result (m-M)$_{LMC}$ = 18.07 $\\pm$ 0.12. ",
        "conclusions": "Our alternate red clump distance modulus (m-M)$_{LMC}$ = 18.36 $\\pm$ 0.17 is $\\approx$0.3 mag longer than that obtained under the assumption that the LMC red clump mimics the Galaxy's.  This is consistent with the ``long'' distance modulus of 18.50 $\\pm$ 0.10 adopted by the Cepheid Key Project, and, contrary to \\cite{uda98} and \\cite{sta98b}, in agreement with the most recent calibrations of the Cepheid period-luminosity relation, which give (m-M)$_{LMC}$ = 18.44 $\\pm$ 0.35 or 18.57 $\\pm$ 0.11 (\\cite{mad98}).  Our value of 18.36 $\\pm$ 0.17 is also in agreement with the shorter distance from RR Lyrae stars of (m-M)$_{LMC}$ = 18.28 $\\pm$ 0.13 (\\cite{lay96}).  Taking account of the stellar population of the SMC pushes that galaxy slightly farther away as well, from (m-M)$_{SMC}$ = 18.56 $\\pm$ 0.09 derived by \\cite{uda98}, to (m-M)$_{SMC}$ = 18.82 $\\pm$ 0.20. The larger value is in good agreement with that from the Cepheid period-luminosity relation, (m-M)$_{SMC}$ = 18.94 $\\pm$ 0.04 (\\cite{lan94}).  The detailed stellar physics of mass-loss, and the relations between Z, Y, and M$_c$ introduce significant uncertainty into the determination of red clump absolute magnitudes, even in the I band.  It is quite likely that these effects may work to bring the Magellanic Clouds to a distance more consistent with the ``short'' (RR Lyrae) distance scale than the ``long'' (Cepheid) scale.   However, the red clump method does not require Magellanic Cloud distances as much as 15\\% smaller than commonly accepted.  Figure 1 shows the menagerie of recent LMC distance determinations with errorbars, and it can be seen that the red clump method gives results consistent with most of the other determinations.  We conclude that the red clump is indeed an extremely useful distance indicator, as described by (\\cite{pac98}, \\cite{sta98}, and \\cite{uda98}). However, like most stellar ``standard candles'', its properties vary with the composition and age of the host galaxy, and the assumption that all red clumps are identical to the {\\sl Hipparcos} red clump is probably incorrect.   Among populations dominated by stars older than $\\approx$6 Gyr, the standard candle approximation should be valid to a high degree for M$_I^0$.  For younger populations, M$_I^0$ is probably brighter than M$_I^0$(local), and the correction may amount to as much as several tenths of a magnitude."
    },
    "9804/astro-ph9804260_arXiv.txt": {
        "abstract": "If the first stars formed soon after decoupling of baryons from the thermal cosmic background radiation (the CBR) the radiation may have been last scattered in a cloudy plasma. We discuss the resulting small-scale anisotropy of the CBR in the limit where the plasma clouds are small compared to the mean distance between clouds along a line of sight. This complements the perturbative analysis valid for mildly nonlinear departures from homogeneity at last scattering. We conclude that reasonable choices for the cloud parameters imply CBR anisotropy consistent with the present experimental limits, in agreement with the perturbative approach. This means the remarkable isotropy of the CBR need not contradict the early small-scale structure formation predicted in some cosmogonies. ",
        "introduction": "\\label{sec-intro}  Cosmogonies in which the baryons were concentrated in clouds at the epoch of last scattering of the thermal cosmic background radiation (the CBR), as in isocurvature models (\\cite{pe94}, \\cite{pe97a}), are in qualitative agreement with the appearance in quasar absorption line spectra of a well-advanced state of structure formation at redshift $z\\sim 5$, but the large amplitude of the density fluctuations may produce significant small-scale anisotropy in the CBR. The analysis of the CBR anisotropy in second order perturbation theory was introduced by \\cite{os86}. Vishniac's (1987) more detailed investigation has been confirmed by \\cite{hu94}, \\cite{do95}, and \\cite{hu97}. The analysis and numerical evaluation are extended in \\cite{pe95} and Persi et al. (1995). The application of perturbation theory may be uncertain if early structure formation produced highly nonlinear departures from homogeneity at the epoch of last scattering, however. To investigate this we have developed a nonperturbative model for scattering in a cloudy distribution of plasma. The expression for the CBR anisotropy $\\delta T/T$ in the strongly cloudy limit bears a close resemblance to the Ostriker-Vishniac relation, and the numerical results for $\\delta T/T$ accordingly are similar.  Our analysis, which extends previous discussions by \\cite{ho89}, \\cite{pe90}, \\cite{ag96} and \\cite{gr98}, assumes the CBR was last scattered by free electrons in clouds with density contrast well above unity, so the mean free distance $t_f$ between intersections of gas clouds along a line of sight is much larger than the typical cloud size $\\dcl$. The simplification offered by this limit is that the details of the matter distribution and motion on scales from $\\dcl$ to $t_f$ are not important, because a line of sight on average samples only one cloud over the distance $t_f$. Thus we can model the clouds as an inhomogeneous random Poisson process. In the process the joint distribution of cloud motions and scattering optical depths as a function of position along the line of sight is determined by the plasma density and velocity fields averaged through a window of width $t_f$. If $t_f$ is larger than the scale of non-linear density fluctuations then the CBR anisotropy $\\delta T/T$ is the sum in quadrature of a perturbative contribution and a shot noise term.  The next section shows the relation between the shot noise term in $\\delta T/T$ and the Ostriker-Vishniac (1986) effect. The model of the clouds as an inhomogeneous random Poisson process is presented in \\S 3, and a simplified treatment of the effect of correlated cloud motions is discussed in \\S 4. In \\S 5 we present numerical examples of the expected CBR anisotropy. ",
        "conclusions": "\\label{discussion}  The simplifying assumption for this analysis is that structure formation is so well advanced at the epoch of last scattering of the CBR by free electrons that the mean distance $t_f$ between intersections of clouds along a line of sight is large compared to the scale of nonlinear mass fluctuations. This allows us to model the clouds as an inhomogeneous random Poisson process determined by the mass density and peculiar velocity fields smoothed through a window of width $t_f$, and it leads to the shot noise contribution to the small-scale CBR anisotropy in equation~(\\ref{eqn-unsaturated}). This approach is motivated by the isocurvature CDM model for structure formation (\\cite{pe97a}), in which structure formation could commence at decoupling at redshift $z\\sim 1000$. A second important motivation has been to complement the usual perturbative analysis of the effect of the nonlinear growth of small-scale structure. The similarity of results from the perturbative (\\cite{psco95}) and nonperturbative approaches leads us to believe we have reliable methods for estimating the effect of early nonlinear structure formation on the CBR anisotropy.  The observations of young galaxies and the intergalactic medium at $z\\sim 3$ indicate a situation intermediate between the perturbative and nonperturbative cases. The damped Lyman-$\\alpha$ systems contain a significant baryon fraction, and the mean distance between intersections of these clouds is large (at $z=3$ it is longer than the Hubble length). There also is a significant baryon fraction in the Lyman-$\\alpha$ forest, and these clouds have a relatively short mean free distance. It is not unreasonable to speculate that the situation at much larger redshifts similarly calls for a combination of the two approaches to the analysis of the angular distribution of the CBR.  The CBR anisotropy produced by the Sunyaev-Zel'dovich (1970) effect of the hot electrons in clusters of galaxies, which certainly is dominated by the shot noise term, offers an important constraint on the epoch of collection of the intracluster plasma. The evidence from the analysis of Persi et al. (1995) is that this constraint does not yet rule out the early structure formation picture. And our conclusion from the numerical examples in \\S 5 is that within presently known observational constraints structure formation could have commenced when the universe was optically thick to scattering of the CBR."
    },
    "9804/astro-ph9804056_arXiv.txt": {
        "abstract": "We present the results of soft X-ray observations of the intermediate-age open cluster IC\\,4651 performed with the ROSAT PSPC. We detected 25 sources. Two are identified with a giant binary and a blue straggler respectively, both belonging to the cluster, and two with probable main-sequence members. Two other cases have ambiguous identification. Of the five binaries known in the cluster, the one detected in X rays is the only one whose period is short enough to maintain fast rotation and therefore strong stellar activity at a high age. The detected blue straggler is probably binary, suggesting that binarity is the key to producing a high level of X-ray emission. It is the third blue straggler detected in X rays. The remaining sources need to be identified through optical follow-up. ",
        "introduction": "Stellar activity depends crucially on the star's rotation rate (Pallavicini et al. 1981), which decreases with age because of magnetic braking (e.g. Skumanich 1972). Thus, open clusters are crucial to distinguish between truly evolutionary effects on stellar activity and effects primarily due to the rotation rate itself. The first X-ray observations of open clusters were carried out with the {\\it Einstein} satellite (Stern et al. 1981, Caillault \\& Helfand 1985, Micela et al. 1988, Schmitt et al. 1990, Micela et al. 1990), but these observations have been carried out in a more systematic way with the ROSAT satellite (Stern et al. 1992; Stauffer et al. 1994; Patten \\& Simon 1993; Randich \\& Schmitt 1995; Randich et al. 1995, 1996a,b). Various clusters of different ages have been studied in order to understand the evolution of stellar activity with age (see Randich 1997 for a review). Before ROSAT, the attention has been concentrated mainly on young clusters (30 to 700 Myr). ROSAT performed the first observations of the old open cluster M\\,67 ($\\sim$5\\,Gyr, Belloni et al. 1993, Belloni et al. in preparation) and NGC\\,188 ($\\sim$9\\,Gyr, Belloni et al., in preparation), leading to the detection of a number of sources. Indication of chromospheric activity has been found for most of the optical candidates (Pasquini \\& Belloni 1994).  In contrast to younger clusters, clusters older than $\\sim$\\,1 Gyr are not expected to contain rapidly-rotating single late-type stars, and therefore strong X-ray sources. However, there are stars older than $\\sim$ 1 Gyr that show rapid rotation: these are members of close binary systems, where tidal interaction prevents the stars from losing angular momentum; well-known examples are the RS CVn binaries. The observations of these clusters have led to the detection of such binaries, and also to a number of peculiar and interesting objects, such as white dwarfs, catalclysmic variables, blue stragglers and wide/eccentric binaries.  In the framework of a project to cover the X-ray observational gap between old and young clusters, we observed the intermediate age open clusters NGC\\,752 ($\\sim$2\\,Gyr, Belloni \\& Verbunt 1996), NGC\\,6940 ($\\sim$1\\,Gyr, Belloni \\& Tagliaferri 1997) and IC\\,4651 ($\\sim$2.5\\,Gyr) with ROSAT PSPC. Moreover, ROSAT HRI observations of the clusters NGC 3680 ($\\sim$2\\,Gyr) and NGC\\,2527 ($\\sim 1$\\,Gyr) were made in 1997 and are currently being analyzed. In the PSPC observation of NGC\\,752, 49 X-ray sources have been detected; seven of them are identified with optical cluster members, four of which are short period binaries, one is a rapid rotator and one is a blue straggler (Belloni \\& Verbunt 1996). In the PSPC observation of NGC\\,6940 18 sources were detected, four of which are identified with members of the cluster with a fifth source a suspected member. In NGC\\,6940, a high fraction of the detected members are binaries: three out of four of the identified members are among the only six binaries known in the cluster. These observations give also evidence for the presence of a saturation level, at which the whole surface of the star is chromospherically active (see Belloni 1998 for a review).  Here we present the results obtained for IC\\,4651. This open cluster has an estimated age of $\\sim 2.5$\\,Gyr and a distance of $\\sim 800 - 900$\\,pc (Eggen 1971, Anthony-Twarog et al. 1988). It has a relatively constant reddening, estimated to be E(B-V)$ = 0.15$ by Eggen (1971) and E(B-V)$ = 0.09$ by Anthony-Twarog \\& Twarog (1987), and an angular size smaller than $\\sim 15'$. The paper is organized as follows: in Sect. 2 we present the PSPC observation and our data analysis, in Sect. 3 we present and discuss the results and in Sect. 4 we compare them with those of other open clusters and discuss the implications. ",
        "conclusions": "The age of IC\\,4651 is about twice that of NGC\\,6940 and similar to that of NGC\\,752. As in the case of NGC\\,752 (Belloni \\& Verbunt 1996), we expect to detect coronal sources only if they are in binary systems; moreover, due to saturation effects one would expect to detect more easily giant stars than main sequence stars (see discussions in Stauffer et al. 1994; Belloni \\& Verbunt 1996; Belloni \\& Tagliaferri 1997). As can be seen from Fig. 1, in IC\\,4651 there are more than a dozen giants. Five of them are found to be binaries by Mermilliod et al. (1995). However, only one of them has a period short enough to sustain a high level of stellar activity, and this is the binary we detect as an active X-ray source. In this scenario, one expects the stars to co-rotate and therefore the orbit to be circular. This binary is slightly eccentric (e=0.1), but its characteristics are compatible with co-rotation (see Verbunt \\& Phinney 1995). This is also the brightest source of the six that we suppose to be at the cluster distance, in agreement with the presence of a saturation level, which is higher for evolved stars (see Belloni \\& Verbunt 1996).  In NGC\\,6940 we had not detected main sequence stars, but only giants. Of these, two are not classified as binaries; one in particular has been extensively studied (Mermilliod \\& Mayor 1989) and is very unlikely to be a binary (see Belloni \\& Tagliaferri 1997). Thus, their detection in the X-ray band is puzzling. On the contrary, in IC~4651 of all known giants only one, a binary, is detected. Two other stars lying on the cluster main sequence are also detected. Given the age of the cluster, we expect them to be binaries. At the bottom end of the cluster main sequence, somewhat offset from it, lie the two possible counterparts of source \\#18. They could be either field stars or weak members of the cluster.  In the error box of source \\#11, which is only 13 arcsec, there are at least seven stars, which look like a sub-cluster. The three brightest of these stars are plotted as crosses in Fig. 1. If these stars are all members of the cluster, then one is a blue straggler (i.e. it lies in a region of the color-magnitude diagram occupied by blue stragglers), while the other two are stars in the process of evolving toward the red giant branch. The possibility cannot be ruled out that more than one of these stars contribute to the detected X-ray emission.  Finally, source \\#14 is identified with a blue straggler, found by Anthony-Twarog \\& Twarog (1987) to be a binary. This is the third blue straggler detected in the X-ray band, the other two being in M67 and NGC\\,752 respectively. Of all blue stragglers known in the old and intermediate-age open clusters that we studied, only three are detected in X rays. All three have measurements that indicate binarity (see Belloni 1998), and once again this is probably the key to X-ray emission. These are not the only binaries know among the blue stragglers in these clusters, showing that binarity does not imply strong stellar activity per se. One of the two stars also has to be of late spectral type in order to have an enhanced dynamo activity, like in Algol systems.  For the sources marked with `OUT'  in Table 1 we have no color information; however, from their magnitudes, most of them should be stars. They lie outside the inner 10 arcmin diameter region of the cluster and are probably non-members. We have already planned follow-up optical observations to determine the physical nature of all detected X-ray sources in the IC\\,4651 field."
    },
    "9804/astro-ph9804146_arXiv.txt": {
        "abstract": "% OH(1720 MHz) masers unaccompanied by 1665/7 MHz line masers have recently been proposed as indicators of the interaction of supernova remnants (SNRs) and molecular clouds.  We present a model for the masing region in which water produced in a C-type shock wave driven into the molecular cloud is dissociated as a result of the X-ray flux from the SNR.  We note that the magnetic field strengths inferred from Zeeman splitting of the 1720 MHz line measure the internal pressure of the supernova remnant.  In addition, we discuss the interaction of Sgr A East, a SNR candidate, with the 50 km/s cloud at the Galactic Centre and present near-infrared observations of \\hmol\\ emission towards the regions where OH(1720 MHz) maser emission is concentrated.  The magnetic field strength obtained from earlier Zeeman measurements is consistent with rough pressure equilibrium between the postshock gas and the X-ray gas filling Sgr A East detected by ASCA. Further, the intensity of the v=1--0 S(1) line of \\hmol\\ is consistent with the shock strength expected to be driven into the molecular gas by this pressure.  The relative intensities of the \\hmol\\ lines in Sgr A East imply mainly collisional excitation. ",
        "introduction": "OH masers have generally been used as a diagnostic for HII regions and evolved stars.  However, a recent study by Frail, Goss \\& Slysh (1994) revealed that the 1720 MHz transition of OH maser emission, when unaccompanied by the 1665 and 1667 MHz OH lines, can be an effective indicator of shock waves interacting with molecular clouds, particularly for supernova remnants (SNRs).  Frail et al.  (1994) detected 26 distinct OH(1720 MHz) maser spots along the interface between the SNR W28 and an adjacent molecular cloud. Evidence of the association between the molecular cloud and W28 comes from the distribution of the molecular material following the eastern edge of the supernova shell (Wootten 1977).  There are more than a dozen Galactic sources and three extragalactic sources in which 1720 MHz OH masers have been found interior to SNR's with adjacent molecular clouds (Yusef-Zadeh, Uchida \\& Roberts 1995; Frail et al. 1996; Yusef-Zadeh et al. 1996; Green et al. 1997; Seaquist, Frayer \\& Frail 1997).  These masers are close both in position and velocity to the interfaces between the remnants and the clouds.  In addition, in several of these sources, CO emission lines reach a maximum in both brightness and linewidth at the interface between remnant and cloud (Wootten 1977).  These are strong observational indications that the shocks are caused by the remnants expanding into their respective adjacent molecular clouds.  Further support is provided by theoretical studies of the pumping of the OH maser lines (Elitzur 1976; Pavlakis \\& Kylafis 1996a,b).  The 1665 and 1667 MHz masers are pumped by far-infrared radiation and are therefore associated with HII regions and evolved stars.  The OH(1720 MHz) maser is collisionally pumped in molecular gas at temperatures and densities between 15-200 K and \\(10^4-10^6\\) cm\\(^{-3}\\), respectively.  Thus in the absence of the 1665/7 MHz transitions, the OH 1720 MHz line presumably traces cooling, shocked gas.  However, shock chemistry predicts that OH is \\emph{not} abundant in the postshock gas as it is rapidly converted to \\water\\ within the shock front.  OH masers adjacent to compact HII regions are produced by photodissociation of \\water\\ (Elitzur \\& de Jong 1978; Hartquist \\& Sternberg 1991; Hartquist et al. 1995), but in that case there is a strong dissociating FUV flux from the star. The dissociating flux is largely absorbed and reradiated in the FIR by grains, providing a sufficient IR background to also pump the 1665/7 MHz transitions (e.g. Pavlakis \\& Kylafis 1996b).  This cannot be the case for the unaccompanied 1720 MHz masers associated with SNR-molecular cloud interactions.  Here we propose that it is the weak X-ray flux from the SNR interior that is ultimately responsible for the dissociation of \\water\\ in the shocked molecular gas.  We note that if the OH maser arises in the postshock gas, then Zeeman measurements determine the magnetic field strength in the postshock gas, and thus measure the pressure within the SNR more directly than other methods.  The observations of the IC 443, W28 and W44 are consistent with this scenario.  Finally, we apply these ideas to the interaction of the Galactic center nonthermal source Sgr A East, which is either a SNR or a multiple-SNR driven bubble, with a molecular cloud at the Galactic center. ",
        "conclusions": "We have proposed a model for the OH 1720 MHz masers unaccompanied by main-line transitions that are associated with SNRs interacting with molecular clouds (Frail et al. 1994, 1996; Green et al. 1997).  A C-type shock wave driven into the adjacent cloud produces water that is subsequently dissociated by the secondary FUV flux produced by the interaction of X-rays from the SNR incident on the molecular cloud. The dissociation only becomes significant once the shocked molecular gas has cooled to about \\( 180 \\u K \\), naturally producing OH at the temperature and abundance required for collisional pumping of the 1720 MHz transition.  Some consistency checks suggest that this model also applies to the OH 1720 MHz masers seen towards Sgr A East.  Firstly, within the framework of the model the magnetic field strengths of roughly 3 mG inferred from Zeeman splitting of the OH masers (Yusef-Zadeh et al. 1996) are measurements of the magnetic pressure in the postshock gas, which dominates the gas pressure behind the shock front.  Equating the magnetic pressure, \\( 4\\ee -7 \\u erg \\percc \\), to \\( \\rho v_s^2 \\), where \\( \\rho \\) is the preshock density and \\( v_s \\) is the shock speed, and adopting a preshock density of \\( \\nh=2 \\ee 4 \\percc \\) (Mezger et al.  1989), we infer \\( v_s \\approx 25\\)--\\(30 \\kms \\).  A C-type shock at this speed produces an intensity in the 1-0 S(1) line of \\( 10^{-4} \\)--\\(10^{-3} \\u erg \\ut s -1 \\ut cm -2 \\) (Draine et al. 1983; Kaufman \\& Neufeld 1996).  The \\emph{measured} intensity, after correcting for extinction assuming that \\( A_K \\approx 2.5 \\), is \\( 1.5\\ee -4 \\u erg \\ut s -1 \\ut cm -2 \\). Second, the inferred postshock pressure is comparable to the pressure of the X-ray emitting gas filling Sgr A East detected with ASCA (Koyama et al. 1996), which has \\( n_e \\approx 6 \\percc \\) and \\( T \\approx 10 \\u keV \\), and a pressure of \\( 2\\ee -7 \\u erg \\percc \\).  Finally, we note that the X-ray luminosity from Sgr A East is \\( 10^{36} \\u erg \\persec \\) (Koyama et al. 1996), providing the required dissociating flux for the production of OH behind the shock front.  Rho (1995) suggests that centrally-peaked thermal X-ray emission from SNRs is characteristic of SNR-molecular cloud interactions.  Green et al. (1997) note that all of the SNRs detected in the OH (1720 MHz) line that have been observed in X-rays are members of this class.  Our model appears to strengthen this link by relying on the X-ray flux from the SNR to produce OH, but in principle a cosmic-ray flux enhanced by a factor of 100 over the solar neighbourhood value has a similar effect on the heating and chemistry of molecular gas (see, e.g.  Maloney et al. 1996).  The high-energy cosmic-ray flux in the W44, W28 and IC 443 SNRs inferred from EGRET observations is roughly two orders of magnitude larger than the solar neighbourhood value (Esposito et al. 1996), so a naive extrapolation down to the cosmic-ray energies important for ionisation suggests that the cosmic-ray and X-ray contributions to the secondary FUV flux may be similar.  In any case, this picture supports the suggestion by Frail et al. (1994) that the presence of the OH 1720 MHz line, and the absence of the 1665/1667 MHz lines provide a clear diagnostic of shocked molecular gas."
    },
    "9804/astro-ph9804236_arXiv.txt": {
        "abstract": "We have measured the counts-in-cells fluctuations of 268 Lyman-break galaxies with spectroscopic redshifts in six 9\\arcm$\\times$ 9\\arcm fields at $z\\sim 3$.  The variance of galaxy counts in cubes of comoving side length 7.7, 11.9, 11.4 $h_{100}^{-1}$ Mpc is $\\sigma_{\\rm gal}^2 \\sim 1.3\\pm0.4$ for $\\Omega_M=1$, $0.2$~open, $0.3$~flat, implying a bias on these scales of $\\sigma_{\\rm gal} / \\sigma_{\\rm mass} = 6.0\\pm1.1$, $1.9\\pm0.4$, $4.0\\pm0.7$. The bias and abundance of Lyman-break galaxies are surprisingly consistent with a simple model of structure formation which assumes only that galaxies form within dark matter halos, that Lyman-break galaxies' rest-UV luminosities are tightly correlated with their dark masses, and that matter fluctuations are Gaussian and have a linear power-spectrum shape at $z\\sim 3$ similar to that determined locally ($\\Gamma\\sim 0.2$).  This conclusion is largely independent of cosmology or spectral normalization $\\sigma_8$. A measurement of the masses of Lyman-break galaxies would in principle distinguish between different cosmological scenarios. ",
        "introduction": "Much of observational cosmology depends upon the assumption that the spatial distribution of galaxies is related in a simple way to the underlying distribution of matter.  At first it was hoped that the galaxy distribution might simply be a Poisson realization of the matter distribution; but as this model became difficult to reconcile with large scale peculiar velocities, the amplitude of microwave background fluctuations, the different clustering strengths of different galaxy types, and theoretical prejudice for $\\Omega_M=1$, cosmologists began to assume an unspecified constant of proportionality $b$ between galaxy and mass fluctuations: $\\delta_{\\rm gal} = b \\delta_{\\rm mass}$.  Though many physical processes could in principle give rise to a relationship of this form (e.g. Dekel \\& Rees 1987), most were poorly understood, and, if invoked, would make it difficult to use galaxy observations to constrain the cosmological mass distribution.  An important exception was gravitational instability.  This is relatively well understood, and if it were dominant in determining where galaxies formed---if galaxies formed within virialized ``halos'' of dark matter, and if the poorly understood physics of star formation, supernova feedback, and so on were important only in determining the properties of galaxies within dark matter halos---then the large scale distribution of galaxies would still be related in a simple way to the underlying distribution of matter; the value of the ``bias parameter'' $b$ in $\\delta_{\\rm gal} = b \\delta_{\\rm mass}$ would be straightforward to calculate (White \\& Rees 1978, Kaiser 1984, Bardeen \\et 1986, Mo \\& White 1996).  Because it maintains a simple relationship between galaxies and mass, agrees with our limited knowledge of the relevant physics, and seems consistent with numerical simulations, this ``dark halo'' model has become increasingly popular, and is now the basis of the modern understanding of galaxy formation.  It is assumed in most analytic treatments, in semi-analytic models, and in numerical simulations which include only gravity; and yet it remains a conjecture that has never been thoroughly tested. One prediction of the dark halo model is that galaxies of a given mass should form first in regions where the density is highest, and since such regions are expected to be strongly clustered (e.g. Kaiser 1984), a natural test is to measure the clustering of galaxies in the young universe.  The Lyman-break technique (e.g., Steidel, Pettini, \\& Hamilton 1995) provides a way to find large numbers of star-forming galaxies at $z\\sim 3$.  Star-forming galaxies have pronounced breaks in their spectra at 912 \\AA\\ (rest) from a combination of absorption by neutral hydrogen in their interstellar media and the intrinsic spectra of massive stars. At $z\\simgt 3$ this ``Lyman break,'' strengthened from additional absorption by hydrogen in the unevolved intergalactic medium, is redshifted sufficiently to be observed with ground-based broad-band photometry.  By taking images through filters that straddle the redshifted Lyman break, and looking for objects that are much fainter in images at wavelengths shortward of the break than longward of the break, one can efficiently separate high-redshift galaxies from the many foreground objects. In our implementation of the technique, we have used deep photometry in the custom $U_n$, $G$, ${\\cal R}$ filter system of Steidel \\& Hamilton (1993) to assemble a sample of over~1300 probable $z\\sim 3$ galaxies, of which more than~400 have been spectroscopically confirmed with the Low-Resolution Imaging Spectrograph (Oke \\et 1995) on the W.~M.~Keck telescopes.  After initial spectroscopy in one 9\\arcm$\\times$ 18\\arcm field, we argued, on the basis of a single large concentration of galaxies, that these $z\\sim 3$ Lyman-break galaxies were much more strongly clustered than the mass, with an inferred bias parameter $b\\equiv \\sigma_{\\rm gal}/\\sigma_{\\rm mass}$ of $b\\simgt 6$, $2$, $4$ for $\\Omega_M=1$, $0.2$~open, and $0.3$~flat (Steidel \\et 1998a). Qualitatively this strong biasing was consistent with the idea that galaxies form first in the (strongly clustered) densest regions of the universe, but there appeared to be quantitative problems. In the dark matter halo model there is an inverse relationship between the abundance and bias of a population of halos, with the rarest, most massive halos being the most strongly clustered (i.e., most ``biased''). As emphasized by Jing \\& Suto (1998), for halos to be as strongly clustered as Lyman-break galaxies, they would have to be very rare indeed. Yet Lyman-break galaxies are not that rare; for $\\Omega_M=1$ their comoving number density to ${\\cal R}=25.5$ is $n\\simgt 8\\times 10^{-3}$ per $h_{100}^{-3}$ Mpc$^3$, comparable to the number density of $L_*$ galaxies today.  As we will see below, in standard ($\\Omega_M=1$, $\\sigma_8=0.6$, $\\Gamma=0.50$) CDM, halos at $z\\sim 3$ with the same abundance as observed Lyman-break galaxies have a bias of $b\\sim 4$, substantially lower than the implied galaxy bias. For $\\Omega_M<1$ the disagreement is less severe, because both the estimated bias and the comoving abundance of observed Lyman-break galaxies are lower.  It appeared then from preliminary analyses that our data were consistent with the dark halo model only for $\\Omega_M<1$; but it was unclear how seriously to take conclusions based on a single feature in a single field. Moreover other authors soon analyzed the overdensity differently and argued that it was consistent with models in which galaxies are significantly less clustered than we claimed, with $b$ low enough to remove the inconsistencies with the abundances for $\\Omega_M=1$ (e.g. Bagla 1997, Governato \\et 1998, Wechsler \\et 1998).  In this paper we present a counts-in-cell analysis of the clustering of 268 Lyman-break galaxies (all with spectroscopic redshifts) in six 9\\arcm$\\times$ 9\\arcm fields. This sample contains four times as many galaxies over an area three times as large as our original analysis.  Since in addition it takes into account all galaxy fluctuations in the data, and not just a single over-density, one might hope it would provide a more definitive measurement of the strength of clustering. ",
        "conclusions": "We have estimated the variance of Lyman-break galaxy counts in cubes of side length 7.7, 11.9, 11.4 $h_{100}^{-1}$ Mpc as $\\sigma_{\\rm gal}^2\\sim 1.3\\pm0.4$ for $\\Omega_M=1$, $0.2$~open, $0.3$~flat. This variance implies that Lyman-break galaxies have a bias $b\\equiv\\sigma_{\\rm gal}/\\sigma_{\\rm mass}$ of $b=6.0\\pm 1.1$, $1.9\\pm 0.4$, $4.0\\pm 0.7$ for the same cosmologies. The bias is in good agreement with a simple model, first proposed by White \\& Rees (1978), in which galaxies form within virialized halos of dark matter.  The agreement of our data with this model depends on cosmology primarily through the shape of the power-spectrum, rather than through the growth rate of matter perturbations as might have been expected.  Given the abundance of Lyman-break galaxies and the locally determined power-spectrum shape, one could have predicted a priori from this model the clustering strength we have observed. The agreement is surprisingly good, for it assumes not only that galaxies form within dark halos---which is plausible enough---but that galaxies UV-bright enough for us to detect reside almost exclusively within the most massive halos. UV luminosity depends so strongly on the age of a starburst and on the importance of dust extinction that one might have expected halo mass to play a comparatively minor role in Lyman-break galaxies' UV luminosities; but this appears not to be the case. The observed abundance and clustering properties of Lyman-break galaxies suggest instead an almost one-to-one correspondence of massive halos to {\\it observable} galaxies, and this implies, for example, that the most massive halos essentially always exhibit star formation at detectable levels (i.e., that the duration of star formation is close to the time interval over which the galaxies in the sample are observed), and that halos only slightly less massive rarely do.  The simple analytic approach adopted in \\S 4 cannot justify more precise statements here; these will be presented elsewhere (Adelberger \\et 1998).  While we have argued that our data can be understood through an appealingly simple model for galaxy formation in which galaxies form within dark-matter halos, the UV luminosity of young galaxies is tightly correlated with their mass, and the power-spectrum of mass fluctuations at $z\\sim 3$ has a shape similar to that determined locally, this does not of course rule out other models.  We look forward to learning how well our data agree with competing scenarios for galaxy formation.  In the meantime, one prediction of the scenario we favor is that fainter samples of Lyman-break galaxies in the same redshift range should exhibit weaker clustering; existing data will allow us to test this observationally (Giavalisco \\et 1998b).  \\bigskip \\bigskip  It is a pleasure to acknowledge several conversations with J. Peacock at the beginning of this project. We are grateful to the many people responsible for building the W.~M.~Keck telescopes and the Low Resolution Imaging Spectrograph. Software by J. Cohen, A. Phillips, and P. Shopbell helped in slit-mask design and alignment. CCS acknowledges support from the U.S. National Science Foundation through grant AST 94-57446, and from the Alfred P. Sloan Foundation. MG has been supported through grant HF-01071.01-94A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555."
    },
    "9804/astro-ph9804285_arXiv.txt": {
        "abstract": "Some form of missing energy may account for the difference between the observed cosmic matter density and the critical density. Two leading candidates are a cosmological constant and quintessence (a time-varying, inhomogenous component with negative pressure). We show that an ideal, full-sky cosmic background anisotropy experiment may not be able to distinguish the two, even when non-linear effects due to gravitational lensing are included.  Due to this ambiguity,  microwave background experiments alone may not determine the matter density or Hubble constant very precisely.  We further show that degeneracy may remain even after considering classical cosmological tests and measurements of large scale structure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804128_arXiv.txt": {
        "abstract": "We analyze RXTE/PCA observations of the black hole candidate and galactic superluminal source GRO~J1655--40 during its recent outburst. We show that during its decay to quiescence, GRO~J1655--40 goes through the high, intermediate, and low state (and that at the beginning of its decay it might have even shown signatures of a very high state), just like other black hole candidates.  This is the first time that such a transition is observed in a galactic superluminal source.  We discuss what are the implications of these results on the hypothesis that the spin of the black hole in superluminal sources is much higher than in other black hole candidates. ",
        "introduction": "Our understanding of low magnetic-field accreting neutron stars and black hole candidates (BHC) has greatly improved in recent years, and a general scenario has been proposed in which source properties are determined by basic parameters, such as accretion rate, compact object mass, and magnetic field (van der Klis \\cite{vanderklis94a}).  One of the keys to this unified scenario has been the successful use of simultaneous timing and energy spectral data to identify the different source states, both in neutron star and black hole systems.  For BHCs four states have been recognized:  (1) the low state (LS; Tananbaum et al.  \\cite{tananbaum72}; Oda et al.  \\cite{oda71}), with a flat (photon index $\\Gamma = 1.5-2$) power law X-ray spectrum and strong ($25-50$\\,\\% rms) band limited noise with a break frequency of $0.03-0.3$~Hz; (2) an intermediate state (IS; M\\'endez \\& van der Klis \\cite{mendez97}; Belloni et al.  \\cite{betal97}) where the $2-10$~keV flux is a factor of $\\sim 4$ higher than in the LS, a soft component is present in the energy spectrum while the power law component is softer than in the LS ($\\Gamma = 2-3$), the power spectrum shows $6-8$~Hz Quasi-Periodic Oscillations (QPO), and the break frequency of the band limited noise is in the range $1-10$~Hz; (3) a high state (HS; Tananbaum et al. \\cite{tananbaum72}; Oda et al.  \\cite{oda77}) where the $2 - 10$~keV flux is an order of magnitude higher than in the LS due to the presence of an ultrasoft X-ray spectral component with sometimes a $\\Gamma = 2-3$ power law tail, and a weak (few \\% rms) flat power law power spectrum; and (4) a very high state (VHS; Miyamoto et al.  \\cite{miyamoto91}) characterized by high $2-20$~keV X-ray luminosity ($2-8$ times higher than in the HS), an ultrasoft X-ray spectral component plus a ($\\Gamma \\sim 2.5$) power law tail, strongly variable $1-15$\\,\\% rms band limited noise with a much higher cut off frequency ($1-10$~Hz) than in the LS, and often $3-10$~Hz QPO.  Apart from the difference in $2-20$~keV luminosity, the timing and spectral properties of the IS are quite similar to those of the VHS; the distinction comes from the fact that they are separated by a HS with very different properties.  So far, the VHS has only been observed in GX~339--4 (Miyamoto et al. \\cite{miyamoto91}) and GS~1124--68 (Ebisawa et al.  \\cite{ebisawa94}).  The order in which black hole transients in their decay, and the persistent BHCs, have been observed to move through these states, and similarities to neutron star states strongly suggest that the accretion rate is highest in the VHS, and decreases through the HS, IS, and LS (Miyamoto et al.  \\cite{miyamoto91}; van der Klis \\cite{vanderklis94a}; M\\'endez \\& van der Klis \\cite{mendez97}; Belloni et al. \\cite{betal97}).  Two remarkable BHCs, GRS~1915+105 and GRO J1655--40, so far stood out from the other ones.  They show relativistic ejections, resembling a small scale analogue to the jets observed in active galactic nuclei and quasars (Mirabel \\& Rodr\\'{\\i}guez \\cite{mr94}), and several other types of unique behavior (Morgan, Remillard, \\& Greiner \\cite{mrg97}; Belloni et al.  \\cite{bmk97b}; Fender et al.  \\cite{fender97}), but until now there was little evidence for any of the canonical black hole states described above.  GRO~J1655--40 was discovered during an outburst in 1994, and since then it has shown irregular outburst activity (e.g., Zhang et al. \\cite{zes97}).  The most recent X-ray event of this source started on 1996 April 25 (Remillard et al.  \\cite{remiau96}), and lasted for about one and a half year.  Dynamical measurements suggest the compact star in GRO~J1655--40 is a black hole, with a mass of $\\sim 7$~M$_{\\odot}$ (Orosz \\& Bailyn \\cite{ob97}, van der Hooft et al.  \\cite{hha98}).  A few weeks after the 1994 X-ray outburst, radio jets were detected expanding at apparent superluminal velocities (Tingay et al. \\cite{tingay95}; Hjellming \\& Rupen \\cite{hr95}).  Here we present an analysis of the properties of GRO~J1655--40 during the decay from its latest X-ray outburst.  We find, for the first time in a superluminal source, that GRO~J1655--40 goes through the HS, IS and LS just like other BHCs. ",
        "conclusions": "Our results show that the BHC and superluminal source GRO~J1655--40 fits in with the general picture of BHC states described in Section 1. During these observations the total $2-10$~keV flux decayed by a factor of $\\simmore 5000$, the power law component in the energy spectrum became harder, and the fractional contribution of the soft component to the total flux decreased by a factor of $\\sim 3$.  Simultaneously, the time variability increased by a factor of $\\simmore 3$, and the power spectrum developed a flat-topped band-limited noise component with a break frequency of $\\sim 0.1 - 1$~Hz.  All these changes, both in the energy and power spectra, indicate that GRO~J1655--40 gradually passed through the HS, IS and LS.  This is the first time that such a transition was observed in a galactic superluminal source, and it provides a unique opportunity to understand these sources in terms of canonical black hole states.  One of the characteristic properties of the band limited noise component of BHC in the LS and IS (and the VHS) is its variable break frequency, which anticorrelates with the rms amplitude measured at the break (Belloni \\& Hasinger \\cite{b&h90}, M\\'endez \\& van der Klis \\cite{mendez97}).  The PDS of August 14 and 18 show the same anticorrelation and fit in with previously observed sources (see M\\'endez \\& van der Klis \\cite{mendez97}), supporting the conclusion that GRO~J1655--40 behaves similarly to other BHCs.  On July 29 and August 3 the PDS are consistent with those of a typical HS (van der Klis \\cite{vanderklis94a}), but on May 28 and July 8 they do not fit the canonical $\\alpha \\sim 1 - 1.5$ single power law shape.  Due to the small rms amplitude that characterizes it, the PDS in the HS was not very well constrained in most previous observations.  It may well be that the ``unusual'' shape of the PDS on May 28 and July 8 (compared to what is expected in the HS) is the rule, but it is only the greater sensitivity of RXTE that allows us to measure it in more detail.  On May 28 the rms amplitude is also larger than was typically observed in the HS, and comparable to that of the ``anomalous'' state reported by Kuulkers, van der Klis, \\& Parmar (\\cite{kvp97}) in 4U~1630--47. Perhaps this is an aspect of VHS behavior.  The May 28 $2-10$ flux is $\\simmore 60$ times higher than in the LS (August 18), and 4 times higher than in the HS (July 29), and both the shape of the power spectrum and the total rms variability are very similar to those observed in GS~1124--68 on February 6 1991, during its VHS (compare the power spectrum in Figure~\\ref{figpow} with that of Figure~1 in Miyamoto et al.  \\cite{miyamoto94}).  This would make GRO~J1655--40 the third BHC (after GX~339--4 and GS~1124--68) to be observed in the VHS.  The application of the DBB plus power law model to BHCs, previously strongly advocated on the basis of Ginga observations of classical BHC (Tanaka \\& Lewin \\cite{tl95}) gains credence by the link provided by our observations between the superluminal and the classical BHC.  In GRS~1915+105, the application of this model leads to a very compelling explanation (Belloni et al.  \\cite{bmk97b}) for the large luminosity excursions in terms of rapid variations of the inner disk radius.  Using the values of $R_{\\rm in}$ and $T_{\\rm in}$ obtained from the spectral fits, and the derived mass of the BHC (Orosz \\& Bailyn \\cite{ob97}; van der Hooft et al.  \\cite{hha98}), we can infer the accretion rate of GRO~J1655--40 during each observation (see Figure~\\ref{figasm}, inset).  Despite the large decrease of $\\dot M$, the DBB flux and the total flux during this transition, the inner radius of the disk remained the same. The fitted values are consistent with a constant radius of 24.8~km.  The same has been observed in GS~1124--68, GS~2000+25, GX~339--4 and LMC~X--3 (Tanaka \\& Lewin \\cite{tl95}, and references therein), although in those cases the flux decreased by less than two orders of magnitude, not so dramatically as in GRO~J1655--40.  Tanaka (\\cite{tan92}) interpreted this lack of change in $R_{\\rm in}$ as a signature of the innermost stable orbit around a black hole.  If we assume that the value of $R_{\\rm in} = 24.8$~km that we obtain from the fits is in fact the radius of the innermost stable orbit, the required spin parameter for a 7~$\\msun$ black hole is $a_{\\rm *} = +0.89$.  Recently, Cui, Zhang, \\& Cheng (\\cite{czc98}) proposed that certain QPO features observed in several BHC are produced by the precession of the accretion disk due to the relativistic dragging of the inertial frame around a spinning black hole.  Using the 300~Hz QPO in GRO~J1655--40 (Remillard et al.  \\cite{remetal97}), and the 67~Hz QPO in GRS~1915+105 (Morgan et al.  \\cite{mrg97}), they conclude that both galactic superluminal sources contain a very rapidly rotating black hole ($a_{\\rm *} \\sim +0.95$).  Similarly, they interpret the 6.7~Hz peak in GS~1124--68 (Belloni et al.  \\cite{betal97}) and the 8~Hz QPO in Cyg~X--1 (Cui et al.  \\cite{cuietal97b}) as the same phenomenon, and conclude that the black holes in these two sources are spinning less rapidly ($a_{\\rm *} \\sim +0.35$, $a_{\\rm *} \\sim +0.48$, respectively). Based on these results, Cui et al.  (\\cite{czc98}) propose that the difference between BHC that are also superluminal jet sources and otherwise ``normal'' BHC is the spin of the black hole.  Both in GS~1124--68 and Cyg~X--1, the QPOs used by Cui et al.  to reach this conclusion were observed during a source transition through the intermediate state, identical to the one in GRO~J1655--40 that we describe here.  Remarkably, on August 14 and 18 the PDS show a QPO feature moving from $\\sim 6.5$~Hz to $\\sim 0.8$~Hz, which is very similar to those observed in GS~1124--68 and Cyg~X--1.  In all three sources the QPO is more prominent at higher energies, and both in GRO~J1655--40 and in Cyg~X--1 the frequency of the QPO is correlated to the break frequency of the band limited noise component.  However, if we, as Cui et al.  (\\cite{czc98}) did in the case of GS~1124--68 and Cyg~X--1, interpret this intermediate-state QPO in GRO~J1655--40 as the precession frequency of the disk, we would derive a spin of the black hole of +0.35 on August 14, and $\\sim -0.1 - + 0.1$ on August 18 (see Cui et al.  \\cite{czc98}, Figure~2), in strong contradiction with the value of +0.89 from the spectral fits.  This argues against the interpretation of the intermediate-state QPO as the precession frequency of the disk.  Of course, these results cannot be used to rule out the possibility that the 300~Hz and 67~Hz features in GRO~J1655--40 and GRS~1915+105 are related to the frame dragging effect, but if they are, then our results provide a strong objection to a similar interpretation of the 6.7~Hz and 8~Hz QPOs in GS~1124--68 and Cyg~X--1.  Either way, we must conclude that the hypothesis that the angular momentum of the black holes in superluminal sources is much higher than that of other BHC is no longer supported by the observed differences in QPO properties between GRO~J1655--40 on the one hand, and Cyg~X--1 and GS~1124--68 on the other.  Finally, we can compare our results for GRO~J1655--40 to those obtained for the other galactic superluminal source, GRS~1915+105.  By means of time-resolved spectral analysis, Belloni et al.  (\\cite{bmk97b}) were able to explain all spectral changes observed in GRS~1915+105 as the effect of thermal-viscous instabilities in the radiation-pressure dominated region of an accretion disk.  Based on the fit results to different parts of the highly variable X-ray light curve, they infer values of $\\dot M$ of $\\sim 1.7-10 \\times 10^{-8} \\msun$~yr$^{-1}$ during the quiescent phase (the range being determined by the uncertainty in the spin of the central black hole), and a factor of two higher during the so-called burst phase.  The accretion rates estimated here for GRO~J1655--40 are about an order of magnitude lower than those of GRS~1915+105 during its quiescent phase.  It is possible that the complicated variability in the X-ray light curve of GRS~1915+105, which is not observed in GRO~J1655--40, is entirely due to the higher accretion rate in that source (although a higher mass, and therefore Eddington luminosity of GRS~1915+105, may compensate this).  A natural consequence of this hypothesis is that, at higher accretion rates GRO~J1655--40 (and other BHCs) should display a similar behavior as the one observed in GRS~1915+105, while on the other hand, at lower accretion rates GRS~1915+105 should undergo a source transition, similar to the one described here for GRO~J1655--40."
    },
    "9804/astro-ph9804291_arXiv.txt": {
        "abstract": "We present deep, pointed $^{12}$CO($J=2-1$) observations of three late-type LSB galaxies. The beam-size was small enough that we could probe different environments (\\HI maximum, \\HI mininum, star forming region) in these galaxies. No CO was found at any of the positions observed.  We argue that the implied lack of molecular gas is real and not caused by conversion factor effects. The virtual absence of a molecular phase may explain the very low star formation rates in these galaxies. ",
        "introduction": "Low Surface Brightness (LSB) galaxies exhibit some of the most extreme properties known for disk galaxies.  This class itself can be split into the group of disk dominated, late-type field galaxies with absolute magnitudes between $M_B \\sim -17$ and $M_B \\sim -19$ and central surface brightnesses $\\mu_0(B) \\simeq 23.5$ mag arcsec$^{-2}$, and the much smaller group of giant LSB galaxies (``Malin-1 cousins'', Sprayberry et al 1995 and Knezek 1993).  Judging from their morphology these latter galaxies have undergone an evolutionary history quite different from the late-type field LSB galaxies which we will consider in this paper.  The extreme gas-richness (McGaugh \\& de Blok 1997) and low metallicities (McGaugh 1994) of the late-type LSB galaxies indicate that they are quite unevolved.  They have low mass surface densities and this has often been suggested as a possible cause for their slow evolution (van der Hulst et al.\\ 1987, McGaugh 1992, de Blok \\& McGaugh 1996).  Detailed investigations of a small sample of LSB galaxies (van der Hulst et al.\\ 1993) show that their gas surface densities in general lie below the critical density needed for star formation, as derived by Kennicutt (1989).  Although this global threshold density should be considered as a boundary condition only (local instabilities may still cause star formation), it nevertheless shows that conditions for star formation in LSB galaxies are not as favourable as in ``normal'' high surface brightness (HSB) galaxies.  This might simply be caused by the low densities, making dynamical timescales much longer, and therefore hampering the collapse of gas complexes into Giant Molecular Clouds (GMCs).  The low metallicity may also make cooling of the Interstellar Medium (ISM) more difficult, delaying the formation of GMCs.  To get a better handle on the properties of the cold, molecular component of the ISM in LSB galaxies, one needs to observe indicators such as the CO molecule.  Previous studies (Schombert et al. 1990) that have tried to detect CO, have not succeeded to rather low limits. This either means that CO does not work as a tracer in LSB galaxies -- implying that large amounts of molecular hydrogen could still exist -- or that LSB galaxies are deficient in their molecular component.  The case of H$_2$-poor galaxies is especially interesting: the conditions that can then be deduced for LSB galaxies, which have obviously formed stars, might help answer questions as: where and how do stars form in an environment poor in molecular gas?  Is a small molecular component, even as an intermediate agent always needed? Are GMCs always needed for star formation?  In this paper we will describe the results and implications of a few very deep pointed observations in the \\hbox{CO(2-1)} line of various galactic environments in LSB galaxies. The higher resolution of this line with respect to the \\hbox{CO(1-0)}-line enabled us to point at different locations within one galaxy.  Positions were selected on the basis of detailed H{\\sc i} (de Blok et al.\\ 1996) and optical imaging (de Blok et al.\\ 1995).  Section~2 describes the observations; in sect.~3 the results are discussed; sect.~4 gives a discussion of the implications; and sect.~5 summarizes the results. We will assume a Hubble constant $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$ in the rest of this paper. ",
        "conclusions": "We have presented deep, pointed $^{12}$CO($J=2-1$) observations of three LSB galaxies. No CO was found at any of the positions observed. This leads to a mean upper limit of the local $M({\\rm H}_2)/M({\\rm H{\\sc i}}) < 0.25$, with individual values for the total $M({\\rm H}_2)/M({\\rm H{\\sc i}})$ ratio reaching less than 6\\% (assuming a Galactic value for the CO to H$_2$ conversion factor $X$).  It is, however, unlikely that the Galactic conversion factor applies to LSB galaxies which have low metallicities and low dust content. In fact the value for $X$ is likely to be $\\sim 4$ times higher. The H$_2$ content would be correspondingly higher.  Our limits then imply that LSB galaxies roughly have (less than) 25\\% of their gass mass in the form of H$_2$.  This is still lower than is found in HSB galaxies.  The conclusion then is that there are no large amounts of H$_2$ hidden in LSB galaxies.  The low star formation rates measured in LSB galaxies can thus be explained by the virtual absence of a molecular component. Star formation in LSB galaxies may thus proceed in a different way than in HSB galaxies.  A detailed comparison between the properties of star forming regions in LSB and HSB galaxies may be a good way to put more constraints on the way stars form in environments that lack a cold component."
    },
    "9804/astro-ph9804014_arXiv.txt": {
        "abstract": "We have imaged the nucleus of M32 at 1600 \\AA$ $ (FUV) and 5500 \\AA$ $ (V) using the Wide-Field/Planetary Camera 2 (WFPC2) aboard HST. We detected the nucleus at 1600 \\AA$ $ using the redleak-free Woods filter on WFPC2. The FUV light profile can be fit with a Gaussian of FWHM 0\\farcs46 (4.6 pixels), but cannot be resolved into individual stars; no UV-bright nuclear structure was detected. The (FUV$-$V) color of the nucleus is 4.9 $\\pm$0.3, consistent with earlier observations. We are unable to confirm any radial variation in (FUV$-$V) within 0$\\farcs$8 of the nucleus; beyond that radius the FUV surface brightness drops below our detection threshhold.  We also performed surface photometry in V and found our results to be in excellent agreement with deconvolved, WFPC1 results.  M32's light profile continues to rise in a nuclear cusp even within 0\\farcs1 of its center. No intermediate-age stellar population is required by evolutionary population synthesis models to reproduce the (FUV$-$V) color of the nucleus, although these data and current models are insufficient to resolve this issue. ",
        "introduction": "At a distance of just 725 kpc, M32 is the nearest ``true elliptical'' (as contrasted to a dwarf elliptical\\footnotemark) galaxy to the Milky Way.  M32's relative proximity allowed its outer regions to be resolved into stars by the first generation of large optical telescopes (\\cite{baa44}), while its crowded inner regions remained unresolved until the advent of the Hubble Space Telescope (see \\cite{gri96}).  The ability to study both the integrated light of M32 and the individual stars contributing to that light has made this galaxy a keystone of efforts to understand the stellar populations of elliptical galaxies beyond the Local Group.  \\footnotetext{see \\cite{dac97} for some illumination of the murky waters of galaxian semantics.}  One of the outstanding puzzles in the study of elliptical galaxies has been that of the ultraviolet excess (UVX); this term refers to the increase in flux shortward of 2000 \\AA$ $ seen in integrated spectra of elliptical galaxies, relative to the amount predicted by the simplest model fits to their optical spectra (\\cite{cod69}, \\cite{cod79}). An important step in identifying the stellar population responsible for the UVX was made by \\cite{bur88}, who identified a strong correlation between the strength of the UVX and the mean metallicity of the galaxy in a sample of 24 quiescent systems.  M32 defined the UV-faint, low-metallicity end of the \\cite{bur88} sample, with an M$_{1550}$ $-$ V-band color of 4.5 $\\pm$0.2 and a mean metallicity of [Fe/H] $\\approx$ $-0.25$ (c.f. the discussion in \\cite{gri96}).  \\cite{bur88} attributed the (1550$-$V) color of M32 to the presence of ``classical'' post-asymptotic giant branch (P-AGB) stars in the galaxy.  The ability of P-AGB stars to produce the (FUV$-$V) color of M32 depends strongly on the age and metallicity distributions of M32's stellar population. However, the very strong degree of central concentration which makes M32 an attractive target for studies of integrated light has hindered progress in the characterization of its stellar populations.  \\cite{dac97} provides an overview of the various strands of evidence for and against the presence of an intermediate-age (few Gyr) population in M32.  Spectral syntheses performed by \\cite{oco80} indicated that the dominant contributors to the present-day optical luminosity of M32 are aged $\\approx$5 Gyr with solar metallicity, but noted that the peak star formation rate might have occurred as long ago as 15 Gyr if the metallicity ranged as low as one-tenth solar.  The existence of a major intermediate-age population in M32 has yet to be proven beyond a reasonable doubt despite the significant body of work dealing with this problem over the past two decades.  Ground-based (e.g., \\cite{fre89}, \\cite{dav92}) and HST (\\cite{gri96}) photometry have both been interpreted to indicate the presence of a significant metallicity spread in M32, in good agreement with recent long-slit spectra (\\cite{har94} 1994). Most investigators seem to agree that significant numbers of stars with metallicities between $-1.5 \\lesssim$ [Fe/H] $\\lesssim 0.0$ are present in M32's stellar mix; of course, the relative numbers of metal-rich and metal-poor stars are poorly constrained due to the usual age-metallicity tradeoff.  \\cite{har94} (1994) discovered the presence of a radial metallicity gradient:  between 15\\arcsec$ $ and 1\\arcmin$ $ from the galaxy's center, the metallicity decreases at the rate $\\Delta$[Fe/H]/$\\Delta$log(r) $= -0.25 \\pm0.07$.  An age gradient has also been reported for M32, as \\cite{gri96} infer from the integrated line indices obtained by \\cite{gon93} that the nucleus of M32 is both several gigayears younger and somewhat more metal-rich than its outer regions.  The deepest optical photometry of resolved stars (\\cite{gri96}) suggests a median age of 8 Gyr for M32 with an average metallicity of [Fe/H] $= -0.25$ for a region 1\\arcmin$ $--2\\arcmin$ $ from the nucleus.  M32 was observed by \\cite{ber95} in the far-ultraviolet using HST's Faint Object Camera (FOC), but there was difficulty in calibrating the absolute throughput of their filter set; M32's FUV color was interpreted as an indication of a stellar population as young as 3 Gyr with slightly sub-solar metallicity. \\cite{bro97} also imaged M32's central region with FOC, using the F175W and F275W filters; their results seemed to indicate a conflict with stellar evolutionary theory, but the interpretation of the data was complicated by the redleak of the FOC ultraviolet filters (see, e.g., the discussion by \\cite{chi97}).  M32 was also observed by the Ultraviolet Imaging Telescope (\\cite{ohl97}); a strong FUV$-$B color gradient was found, demonstrating that the the galaxy's UVX is {\\it weakest} in the center and increases with radius.  We observed M32 as part of a WFPC2 GTO program to study the hot, luminous post-AGB populations and the UV upturn phenomenon in nearby early-type galaxies (\\cite{tra94b}). The use of the F160BW filter with WFPC2 provides a unique constraint on the FUV properties of M32, coupling HST's high resolution (angular scales $\\approx$ 0\\farcs1 pix$^{-1}$) with a zero-redleak (optical transmission $\\ll$ 1\\% T$_{max}$) wide-band filter that permits an accurate determination of the (FUV$-$V) color (c.f. \\cite{wat94}, \\cite{jon96}). ",
        "conclusions": "It seems impossible to reproduce the integrated spectrum of M32 with a single-age, single-metallicity ``simple stellar population'' (SSP) (\\cite{har94} 1994). However, SSPs remain a useful tool for interpreting the (160$-$555) color, because they provide the tools with which to construct more realistic models.  For SSPs, the (FUV$-$V) color is predicted to be a sensitive function of both age and metallicity, and hence a valuable observation for the derivation of star formation histories (\\cite{bre94}).  The exact dependencies of (FUV$-$V) on stellar age and metallicity are highly uncertain, and this limits its predictive power. As discussed in detail by \\cite{yi97}, the age predictions depend critically on poorly modeled processes such as stellar mass-loss. Therefore, the uncertainties in our interpretation are dominated by systematic model effects rather than observational errors.  In order to provide an interpretational starting point, we compared the observed (160$-$555) color of the nucleus of M32 to the SSP predictions of \\cite{tan96}. Table 1 gives the ages of SSPs which are predicted to exhibit (160$-$555) = 4.9 $\\pm$0.3, for a range of metallicities.  In Table 1, $t_y$ indicates the acceptable age range for models in which the FUV light derives from main-sequence stars, while $t_o$ gives the ages for which evolved stars provide the FUV light. The errorbars in Table 1 are {\\it internal, random, observational} errors only, and merely reflect the propagation of photometric error into the model predictions. A thorough discussion of the systematic errors in these models is beyond the scope of this paper (c.f., \\cite{yi98}), but uncertainties in the physics of stellar-mass loss and helium-enrichment limit the absolute precision of such models to several billion years (particularly the older models, which most concern us here). The ages in Table 1 are therefore most useful as relative measures, so that M32's stellar populations may be compared with those of other galaxies using a consistent set of models. As absolute measures, they are subject to extreme uncertainty.  In these models, the ratio of FUV to optical light is not a monotonically decreasing function of age: early in a galaxy's lifetime, (FUV$-$V) increases (gets redder) as the hot upper main-sequence stars die off.  Eventually, as the total luminosity of the SSP continues to fade with age, the FUV contribution of the short-lived P-AGB stars begins to become important, and (FUV$-$V) becomes bluer again.  At sufficiently high metallicities (Z $>$ 2Z$_{\\sun}$), these models predict a point at which low-mass helium-burning stars avoid the AGB in favor of long-lived UV-bright phases, causing a sharp increase in the FUV flux as this behavior sets in at ages $\\approx$5--10 Gyr.  The (160$-$555) color of M32 can be matched by either young or old SSPs.  We rule out the young (0.5--2 Gyr) class of models for two reasons: first, these ages are much younger than even the youngest intermediate-age populations inferred from optical spectroscopy and photometry of M32; and second, models of P-AGB stars (\\cite{vas94}) for such young ages are predicted to have $m_{160}$ $\\approx$ 18.5 at the distance of M32; only one such P-AGB star would be required to account for the FUV light of M32's nucleus, and would be seen as a point source in our images.  The remaining models indicate an SSP age of 12--13 Gyr for M32's nucleus for metallicities less than about 5 times solar.  In this type of SSP, the FUV light is produced by classical P-AGB stars. In the super metal-rich case the age of the SSP drops to $\\approx$6 Gyr due to the production of hot AGB-manqu\\'{e} stars via metal-enhanced mass loss on the RGB (\\cite{tan96}, \\cite{dor95}, \\cite{gre90}). It is important to remember that these age values are subject to large, model-driven systematic errors and should not be considered as definitive determinations.  M32 is not a likely candidate to harbor AGB-manqu\\'{e} stars, as the required metallicity is too high by a factor of two or more.  The best matching SSPs of reasonable metallicities are then 12--13 Gyr old; this is consistent with the results of \\cite{gri96} indicating that half the optical light of M32 comes from stars older than $\\approx$8 Gyr.  Note that M32 is {\\it not} an SSP; any intermediate-age population that contribues to the optical light is predicted to be FUV-faint and thus pushes the required age of the FUV-bright population to even older ages in order to reproduce the integrated (160$-$555) color.  We do not have enough radial coverage in F160BW to permit comparison with the UIT result that M32's FUV light falls off more slowly with radius than its optical light (\\cite{ohl97}).  No gradient in (160$-$555) is seen with the central 0\\farcs8 of M32's nucleus.  If the UIT result is interpreted in the context of SSPs, it would suggest that either M32's metallicity is increasing with radius (contrary to previous results of \\cite{har94} 1994), or that the galaxy's mean age is increasing with radius.  This type of radial age gradient has also been suggested to exist in the dwarf elliptical galaxy NGC 147 (\\cite{han97}) and in the Andromeda I dwarf spheroidal (\\cite{dac96}); it may turn out to be a general feature of small, spheroidal galaxies."
    },
    "9804/astro-ph9804222_arXiv.txt": {
        "abstract": "The bispectrum of the microwave background sky is a possible discriminator between inflationary and defect models of structure formation in the Universe. The bispectrum, which is the analogue of the temperature 3-point correlation function in harmonic space, is zero for most inflationary models, but non-zero for non-gaussian models.  The expected departures from zero are small,  and easily masked by noise, so it is important to be able to estimate the bispectrum coefficients as accurately as possible, and to know the errors and correlations between the estimates so they may be used in combination as a diagnostic to rule out non-gaussian models. This paper presents a method for estimating in an unbiased way the bispectrum from a microwave background map in the near-gaussian limit. The method is optimal, in the sense that no other method can have smaller error bars,  and in addition, the covariances between the bispectrum estimates are calculated explicitly. The method deals automatically with partial sky coverage and arbitrary noise correlations without modification. A preliminary application to the Cosmic Background Explorer 4-year dataset shows no evidence for non-gaussian behaviour. ",
        "introduction": "The question of whether or not the microwave background sky is well-approximated by a gaussian random field is important for distinguishing inflationary and defect models of the early Universe.    Inflationary models predict immeasurably small non-gaussian components, arising from gravity waves \\cite{BMS97} and the Rees-Sciama effect (\\pcite{MGLM95},\\pcite{MSS95}). The difficulty which besets such tests is that the predicted departures from gaussian behaviour for most defect models are small (e.g. \\pcite{Luo94}), and are correspondingly difficult to detect in the presence of noise, which may be instrumental or cosmic variance.  This makes it extremely important to be able to calculate the statistical properties of the non-gaussian discriminants.  In particular, it will probably be necessary to combine the results of a large number of estimates, to obtain a statistically significant departure from gaussianity (if it exists), or to make a convincing case that the sky is indeed gaussian.  Ironically, the best evidence for an inflationary model may well come not from a specific test for non-gaussianity, but rather from the power spectrum.   For inflationary models, the power spectrum of temperature fluctuations is predicted to have a reasonable amount of structure in it, with multiple acoustic oscillation peaks which should be measurable by future satellites such as MAP and Planck (\\pcite{Jungman}, \\pcite{PhaseA}). Should such structures be found, and found to agree with the inflationary model predictions within the errors, the case against any form of defect model would be strong, and even the present knowledge of the power spectrum appears already to rule out many defect models (\\pcite{PST97}).   In this case, the absence of non-gaussian signatures would be a useful confirmation.  However, if the power spectrum turns out not to be well-fitted by inflationary models, the question of the gaussian nature or otherwise of the fluctuations becomes correspondingly more important. There are many ways to approach the problem of determining whether fluctuations are gaussian, in large-scale structure and in the CMB. These include the 3-point function (e.g. \\pcite{Hinshaw94}, \\pcite{FRS93}, \\pcite{LS93}, \\pcite{GLMM94}), the genus and Euler-Poincar\\'e statistic (\\pcite{Coles88}, \\pcite{Gott90}, \\pcite{Luo94b}, \\pcite{Smoot94}),  peak statistics (\\pcite{BE87}, \\pcite{Kogut95}, \\pcite{Kogut96}) and studies of tensor modes in the CMB \\cite{CCT94}. The approach we take here is to investigate the bispectrum, for which some studies have been made in the large-scale structure literature (\\pcite{HBCJ95}, \\pcite{MVH97}, \\pcite{VHM98}) and in the CMB \\cite{Luo94}.   There may be sharper tools for detecting {\\em specific} non-gaussian models, but the rationale for this approach is that the bispectrum offers a {\\em generic} test for non-gaussian models, in the following sense:  a general field will have non-zero connected $n$-point functions at all orders,  and the bispectrum is the lowest statistic (with $n>1$) for which a gaussian field has zero expectation value.  The principal advantages of the approach detailed in this paper are that it deals automatically with masked regions of the sky and correlated noise, and that the estimates of the bispectrum coefficients come with error bars and covariances between the errors.   This last point is particularly important when one realises that a single bispectrum coefficient estimate is unlikely to rule out a model, because the cosmic variance is often larger than the expected signal, and so one is going to need many coefficients in practice.  A final point is that, for gaussian fluctuations, the estimator below cannot be improved, in the sense of having a smaller error bar. ",
        "conclusions": "In this paper, we see that it is possible to construct an estimator for the bispectrum which is lossless, in the sense that it contains as much information on the bispectrum as the entire map.  It is also unbiased, and the covariance properties of the estimators are calculable.  The estimator involves one approximation -- that the departures from gaussianity are small.  In this limit, there is no other method which will lead to smaller error bars.  The fact that the covariance properties are known is important in practical cases, because the bispectrum may well be small in comparison to cosmic variance, so a single estimate is unlikely to be sufficient to rule out many non-gaussian models.  Many estimates (with different $\\ell_1\\ell_2\\ell_3 m_1 m_2 m_3$) would be required.   As a test for non-gaussianity, this method has the advantage that confidence levels can be computed analytically, without recourse to Monte Carlo simulation.  The notion of an optimal method is defined rather precisely in terms of information content and bias, but the issue of whether a method is good or not is wider than this.   There is no doubt that the number of computations required to do this analysis is very large, dominated for reasonable pixel counts by computation of the $R$ coefficients.  This can be aided by precomputing Clebsch-Gordan coefficients \\cite{CG} and using a packed-storage algorithm, and by using parallel computers, for which this problem is ideal. However, it is clear that it will not be possible to deal directly with the entire dataset from Planck without some form of pre-compression, perhaps along the lines above for the low-$\\ell$ modes.  For high-$\\ell$ modes, subdivision of the sky into essentially independent datasets may be required.  This is an inherent problem for high-order statistics, but it may also be required even for the power spectrum.  A further point to note is that the method deals automatically with sky coverage which is not complete, and arbitrary noise correlations. These are the two crucial tests of any method, since they will be a feature of future high-quality experiments such as Planck and MAP. Three-point function estimators have traditionally not computed the covariances directly, but rely on simulation tests to decide significance (e.g. \\pcite{Kogut96}).  This is the disadvantage of many methods (e.g. genus, extrema correlation functions) for which the evaluation of errors may be difficult analytically.  Against this one has, of course, to balance speed advantages.  A preliminary application to the COBE 4-year data shows large-scale bispectrum coefficients consistent with zero.  However, the errors are sufficiently large that these coefficients do not rule out non-gaussian models with confidence."
    },
    "9804/astro-ph9804109_arXiv.txt": {
        "abstract": "We report on further observations of the luminous halo of NGC\\,5907. New V, I and B deep photometry confirms the existence of an extended stellar halo redder than the disk. Our data are consistent with a faint halo, or very thick disk, composed of a metal-rich old stellar population. We propose that it could be the remnant of a merged small elliptical, and we support our hypothesis with N-body simulations. \\keywords       {galaxies: NGC\\,5907                     - galaxies: interactions                  - galaxies: kinematics and dynamics       - galaxies: photometry                    - galaxies: abundances                   - galaxies: halos} ",
        "introduction": "Sackett et al. (\\cite{Sackett}, hereafter SMHB) have found that the edge--on spiral galaxy NGC\\,5907 is surrounded by a faint halo in the optical R band. The existence of this halo has been confirmed by several authors at other wavelengths. Lequeux et al. (\\cite{PaperI}, hereafter Paper I), have observed this halo in two other bands, V and I, and suggest that it becomes redder at increasing distances from the plane of the galaxy. Rudy et al. (\\cite{Rudy}) and James \\& Casali (\\cite{James}) have also detected the halo in J and K; the combination of their brightnesses with those of SMHB and of Paper I gives results difficult to understand. The halo of NGC\\,5907 is clearly stellar in origin but its nature is still unknown. Its relatively red V-I color implies either a normal, metal-rich, old stellar population, or if the stars are metal-poor an initial mass function favoring extremely low masses (see discussion in Paper I). The observations being very difficult because of the faintness of the halo, we decided to perform new observations, described in Sect. 2 and 3. To interpret them, Sect. 4 proposes a scenario in which NGC 5907 has encountered and cannibalized a small elliptical galaxy about 2 Gyr ago. ",
        "conclusions": "In this paper, we reported on new observations made with the CFH telescope which confirm entirely the conclusions reached in Paper I: the luminous halo around NGC\\, 5907 is redder than the galaxy and has the colors of a metal--rich old stellar population. However discrepancies between the present observations and others in the literature have to be understood, which require further observations. We have shown by a N--body simulation that this red halo could result from cannibalism of a low--mass elliptical (a few 10$^9$ M$_\\odot$) after a slow encounter with NGC\\, 5907.   \\begin{figure} \\psfig{figure=Ab161_f5.ps,bbllx=5cm,bblly=8cm,bburx=11cm,bbury=205mm,width=8cm,angle=-90} \\caption[]{  Final plot of the visible stars, when the two galaxies have merged (here at 1900 Myr). The system is rotated to have its mean angular momentum vector along the vertical axis.} \\end{figure}   How frequent is this phenomenon? Nearby spiral galaxies like M 31 often have elliptical companions which might eventually merge with them. Our Galaxy might have cannibalized a satellite which would have produced the thick disk as the result of stellar dynamical heating, the process evidenced by our simulation (Robin et al. \\cite{Robin}). To get statistical information on this capture process, one should conduct a systematic investigation of luminous halos, including measurements of their colors, around a well--defined sample of edge--on spiral galaxies.  \\begin{figure} \\psfig{figure=Ab161_f6.ps,bbllx=0cm,bblly=35mm,bburx=18cm,bbury=18cm,width=8cm} \\caption[]{ Distribution perpendicular to the plane, at a distance of 6kpc from the center, of the stellar component of NGC 5907 (dash), the companion (dots), the total visible mass (full line), and dark halo (dot-dash), at t=1900 Myr.} \\end{figure}"
    },
    "9804/astro-ph9804279_arXiv.txt": {
        "abstract": "In this paper, a meteoric plasma is analyzed from a physical viewpoint, with particular emphasis on its interaction with radio waves.  The attention is drawn to some macroscopic characteristics of a meteoric plasma and it is shown that the electron--ion collision frequency is not negligible, as commonly thought. ",
        "introduction": "When meteoroids enter the Earth's atmosphere, they create long and narrow trails of ionized gas, which can scatter radio waves.  The meteor radio echo theory finds its roots in studies of the ionosphere made toward the end of the twenties (Skellet \\cite{SKELLET}).  But only after the Second World War, and the development of military radar, the correlation between radio echoes and meteor trails became clear (Hey \\& Stewart \\cite{HEY}). The first experiments explicitly devoted to meteor studies were carried out by Pierce (\\cite{PIERCE}), who observed Draconids during the night of 9 to 10 October 1946.  The first theories on the interaction of radio waves with meteors were due to Lovell \\& Clegg (\\cite{LOVELL}), Kaiser \\& Closs (\\cite{KAISER}), Herlofson (\\cite{HERLOFSON}).  Thereafter, between 1950 and 1960, a lot of efforts in this field were undertaken, but after 1960 the interest quickly decayed.  The state of knowledge was well exposed in a classical paper by Sugar (\\cite{SUGAR}) and the book by McKinley (\\cite{MCKINLEY}).  Toward the end of the eighties, the advent of digital technology renewed the interest in forward--scattering as a useful tool for communication channels over the horizon (see Weitzen \\& Ralston \\cite{WEITZEN}).  Today, meteor radars are widely used, even by amateur astronomers, because of low cost (Jenniskens et al. \\cite{JENNISKENS}, Yrj\\\"{o}l\\\"{a} \\& Jenniskens \\cite{YRJOLA}).  Actually, the two main theories for radio echoes are due to Poulter \\& Baggaley (\\cite{POULTER}), who deal with back--scattering, and Jones \\& Jones (\\cite{JONES2}, \\cite{JONES8}, \\cite{JONES3}), who deal with forward--scattering.  We can note that both theories \\emph{do not consider the meteor as a plasma}, but as a simply ionized gas, with a negligible collision frequency. This is an assumption quite common in work on radio echoes from meteors, except for Herlofson (\\cite{HERLOFSON}), who first considered the meteor as a plasma.  Concerning the difference between back--scattering and forward--scattering, it is worth noting that the two systems use different types of radio waves.  A forward--scattering radar uses a continuous sine wave, while a back--scattering one uses a pulsed wave. This influences the mathematical approach to the problem, but also physical theories.  The pulse shape of the back--scattering radar can be represented by a sum of several components with different frequencies. It follows that, during the propagation, the various components tend to change phase with respect to one another, which leads to a change in shape of the pulse.  The dispersion relation thus depends on the frequency and can be expressed in a Taylor series.  A detailed treatment of the pulse propagation in dielectrics is not the subject of this paper: the interested reader can find useful mathematical tools in, for example, Oughstun \\& Sherman (\\cite{PULSE}).  Here we want only to underline that we cannot simply consider the signal emitted by a back--scattering radar as sinusoidal: thus, the extension of the back--scattering echo theory to forward--scattering is not correct.  The purpose of this paper is to settle some basic concepts in meteor physics and we leave a ``cook--book'' approach to other articles. The interaction of sine waves (for the sake of simplicity) in the radio frequency range with meteoric plasma is investigated. We will see that the common assumption about meteors as a collisionless ionized gas does not have any physical ground. ",
        "conclusions": "In this paper, we settled some basic concepts in meteor physics.  We have dealt with the interaction of sine waves, in the radio frequency range, with meteoric plasma. Attention is drawn to some macroscopic characteristics of a meteoric plasma and it is shown that the electron--ion collision frequency is not negligible, as commonly thought.  It is possible to define two meteor classes (overdense and underdense) according to whether the plasma frequency is higher or lower than the radar frequency respectively.  Overdense meteors are divided into two sub--classes (I and II), depending on the ratio between collision frequency and radar frequency.  Taking into account that the meteoric plasma does not have a definite boundary, a simple formula for the calculations of radio--wave attenuation after the forward--scattering is also presented.  This formula allows us to calculate the meteor height, without taking into account the diffusion coefficient.  Further questions can be put: in this work, potassium ion is considered, but other studies can be carried out in order to know the impact of other alkaline and alkaline--earth metals, such as Na, Ca and Mg."
    },
    "9804/astro-ph9804086_arXiv.txt": {
        "abstract": "We  estimate  the evolution of the  contribution of galaxies to the  cosmic background flux at $912 \\rm \\AA$ by means  of a semi-analytic model of galaxy formation and evolution. Such a  modelling has been quite successful in reproducing the optical properties of  galaxies. We assume hereafter the high--redshift damped Lyman-$\\alpha$ (DLA) systems to be the progenitors of present day galaxies, and we design a series  of models which are consistent with the evolution of cosmic comoving emissivities in the  available near infrared (NIR), optical, ultraviolet (UV), and far infrared (FIR) bands along with the evolution of the neutral hydrogen content and average metallicity of damped Lyman--$\\alpha$ systems (DLA).    We  use these  models  to compute  the galactic contribution to the Lyman-limit emissivity and background   flux for $0 \\simeq z \\le 4$. We take into account the absorption of Lyman-limit photons by HI and dust in the interstellar medium (ISM) of the galaxies. We find that the background Lyman-limit  flux due to galaxies  might dominate  (or be comparable to) the contribution from quasars at almost all redshifts if the absorption by HI in the ISM is neglected. The ISM HI absorption results in  a severe diminishing of this flux---by almost three orders of magnitude at high redshifts to between one and two orders at $z \\simeq 0$. Though the resulting galaxy  flux is completely negligible at high redshifts, it is  comparable to the quasar flux at $z \\simeq 0$. ",
        "introduction": "The study of the evolution of the ionization  state of the diffuse intergalactic medium (IGM) is crucial to understanding  the formation of  structures in the universe. An important input in this study is the intensity and evolution of the background  hydrogen-ionizing flux. The universe is observed to be highly ionized at high redshifts (Gunn \\& Peterson 1965, Giallongo {\\it et al. \\/} 1994) and it is widely believed that the main mechanism of ionization is  photoionization by  photons at and below the Lyman-limit. This view is borne out by estimates of  background hydrogen-ionizing flux as inferred by the 'proximity effect'  which fixes the value of this flux,  $J_{\\rm ion } \\simeq 10^{-21 \\pm 0.5}\\, \\rm erg \\, s^{-1} \\, cm^{-2} \\, Hz^{-1}\\, sr^{-1}$ (Bajtlik, Duncan  \\& Ostriker 1988; Bechtold 1994;  Giallongo {\\it et al. \\/} 1996; Cooke {\\it et al. \\/}  1997) for $2 \\la z \\la 4.5$. It is not clear whether the observed quasars can provide sufficient contribution to   the  flux implied by this effect at $z \\ge 3$ (Miralda-Escud\\'e  \\& Ostriker 1990; Giroux  \\& Shapiro 1996; Haardt \\& Madau 1996; Cooke {\\it et al. \\/} 1997). In light of this fact, it was speculated that high-redshift star-forming galaxies could be the  dominant contributer to this background flux (Songaila, Cowie \\& Lilly 1990). At smaller redshifts, there is a greater uncertainty in the level  of this  flux (Kulkarni \\& Fall 1993; Maloney 1993; Vogel {\\it et al. \\/} 1995; Reynolds {\\it et al. \\/} 1995); however, a tight upper limit of $J_{\\rm ion} < 8 \\times 10^{-23} \\, \\rm erg \\, s^{-1} \\, cm^{-2} \\, Hz^{-1} \\, sr^{-1}$ at $z \\simeq 0$ exists from H$\\alpha$ observations (Vogel et al. 1995). If this flux is entirely due  to  the  background hydrogen-ionizing  photons then the contribution of quasars falls  short of the observed value by nearly an order of magnitude at $z \\simeq 0$ (Madau 1992; Haardt \\& Madau 1996).    Recent years have seen tremendous progress in understanding the formation and evolution of galaxies at moderate and high redshifts. Prominent observations which made it possible  are the probing of the evolution of galaxies with $z \\le 1$ by the  Canada-France  Redshift Survey (CFRS) (Lilly {\\it et al. \\/} 1995), the discovery of high-redshift galaxies using the UV drop-out technique (Steidel {\\it et al. \\/} 1996)  in the Hubble  Deep Field (HDF) (Madau {\\it et al. \\/}  1996; Sawicki {\\it  et al. \\/} 1997; Connolly {\\it et al. \\/}  1997), and a better understanding of the evolution of the neutral hydrogen content and average metallicity of DLA systems from $z \\simeq 0$ to $z \\simeq 4$ (Lanzetta {\\it et al. \\/} 1995, Pettini {\\it et al. \\/} 1997). Several theoretical attempts have been made to explain these observations (Madau 1996, 1997; Fall {\\it et al. \\/} 1996; Frenk {\\it  et al. \\/} 1996; Guiderdoni {\\it et al. \\/} 1997b, hereafter GHBM). These observations give valuable clues  about the history of star formation and the comoving emissivity in several  wavebands in the universe upto $z \\simeq 4$ and  have provided a unique opportunity to model and study the contribution of star-forming galaxies to the background hydrogen-ionizing flux (for recent attempts to estimate the hydrogen-ionizing flux from star-forming galaxies for  $z \\le 1.5$ see Giallongo {\\it et al. \\/} (1997); Deharveng {\\it  et al. \\/} (1997)).  The CFRS  observations   give the  evolution of cosmic emissivity  at $\\lambda \\simeq  2800 \\, \\rm \\AA$ (in the rest frame of the emitter) for $z \\leq  1$ while the  HDF observations estimate  the emissivity at $\\lambda \\simeq 1620 \\, \\rm \\AA$ for $2.5 \\leq z \\leq 4$. Our aim in this paper is to extrapolate the results from these important UV wavelengths   to calculate  the emissivity  (and background flux) at the Lyman-limit wavelength. The hydrogen-ionizing photons are emitted by  massive stars with lifetimes $\\simeq  10^7 \\, \\rm years$,  so their flux is not sensitive to the history of galaxy  evolution, which has much larger time scales. Conversely, one can view it as a direct measurement of the instantaneous SFR. The conversion of emissivity from the wavelengths probed by HDF and CFRS to $912 \\, \\rm \\AA$ requires knowledge of three factors: \\begin{itemize} \\item  the absorption of photons with $\\lambda \\le 912 \\, \\rm \\AA$ within the stellar atmosphere. \\item the amount of dust and the spectral index of dust absorption coefficient in the interstellar medium. \\item the amount of neutral hydrogen in the interstellar medium. \\end{itemize} The first factor is fixed by stellar atmosphere  models and depends  mainly on the  choice of  IMF which is guided by our ability to reproduce observed colours of galaxies.   In extrapolating the amount of dust absorption from the observed value at $1620 \\, \\rm  \\AA$ (or $2800 \\, \\rm \\AA$) to the Lyman-limit wavelength, it is important to know the amount of dust in the galaxies as accurately as possible, because $A_{912}/A_{V} \\simeq 5$ as compared to $A_{1600}/A_{V} \\simeq 2.5$. Therefore a small uncertainty in the dust extinction at UV  wavelengths translates into a much larger error in the extinction at the Lyman-limit. We can use FIR data available to fix the amount of dust at $ z \\simeq 0$. At higher redshift, we can fix it by matching the observed  average metallicity of DLA absorbers  (Pettini {\\it et al. \\/}  1997) and by reproducing the cosmic infrared background (CIB) reported by Puget {\\it et al. \\/} (1996).   The most crucial factor in determining the absorption  of hydrogen-ionizing photons  in the ISM is the amount of HI and its distribution in  galaxies. We can model this by using the available data on DLA absorbers  which give the evolution of the  average HI content of these clouds for $0 \\leq z \\leq 4$ (Lanzetta {\\it et al. \\/} 1995). In our study we have to assume that   DLA absorbers belong to  the same population of galaxies as the ones that are observed in emission in the  CFRS and HDF. No clear consensus has been reached on this point yet, but recent observations may support the explanation that the DLA absorbers are either  thick rotating disks or protogalactic clumps  (Haenelt {\\it et al.} 1997;  Prochaska \\& Wolfe 1997). Thereafter We assume the HI, dust, and  stars to be homogeneously distributed in the ISM. This assumption gives the maximum possible HI absorption and allows us to obtain a lower limit on the fraction of Lyman-limit photons escaping from galaxies. We shall also briefly discuss the implications of the clumpiness of  HI.    We use a semi-analytic model of galaxy formation to study the evolution of background hydrogen-ionizing flux for $0 \\simeq z \\la 4$ (for details see, e.g.  GHMB). In \\S 2, we describe the relevant features this model and its  various assumptions.  The basic method of obtaining the background ionizing flux given the comoving emissivity  and various sources of intergalactic absorption is also briefly described in this section. Our results are detailed in \\S 3.  In \\S 4 we discuss various uncertainties and possible improvements in our estimates and summarize our conclusions in \\S 5. ",
        "conclusions": "We estimated the contribution of star-forming galaxies to the background Lyman-limit flux taking  into account the HI and dust absorption  in the ISM of individual galaxies in a self-consistent way with the cosmic star-formation  history. We assumed  that DLA systems correspond to star-forming regions at high-redshift. We conclude that while star-forming galaxies are unlikely to dominate the background hydrogen-ionizing flux at high redshift, they are most likely to do so in the present universe. The current uncertainties of modelling do not allow us to calculate the redshift of cross-over from quasar-dominated to galaxy-dominated background flux. As already discussed, a good discriminator between these  two sources  of ionizing-background  is the softness of their spectra. Future high-resolution observations of metals (e.g. Carbon) in their various ionization states in  the low-column density Lyman-$\\alpha$ clouds  for $0 \\simeq z  \\le 4$ might enable one to find  the transition from quasar-dominated  to galaxy-dominated background flux."
    },
    "9804/astro-ph9804203_arXiv.txt": {
        "abstract": "\\noindent Deep $UBRI$ images of a $15'\\times15'$ field in the outer spheroid of the M31 galaxy obtained using the Kitt Peak National Observatory 4-meter telescope, and $I$ band images obtained with the Keck 10-meter telescope and Low Resolution Imaging Spectrograph, are used to isolate a sample of candidate red giant branch stars located at a projected radial distance of 19~kpc along the minor axis.  These stars are distinguished from the more numerous distant field galaxies on the basis of broadband $U-B$, $B-R$, and $R-I$ colors and image morphology: We isolate objects whose colors are consistent with the long, but relatively narrow locus occupied by red giants in $UBRI$ color space (as observed in Galactic globular cluster giants and predicted by models spanning a wide range of metallicities and ages), and those whose angular sizes are consistent with the stellar point spread function, $\\rm FWHM=0\\farcs6$--$0\\farcs9$ (Keck) and $\\rm FWHM=0\\farcs9$--$1\\farcs5$ (Kitt Peak).  We carry out the same analysis of data on a comparison field with a similar Galactic latitude to the M31 halo field.  The color-magnitude diagram of objects in the comparison field is well described by a superposition of foreground Galactic dwarf stars (in keeping with a standard empirical model of the Galaxy) against a backdrop of contaminating faint blue field galaxies ($I\\ga21$, $B-I\\sim1\\>$--$\\>$2.5), while the M31 halo field contains a clear excess of faint red objects ($I\\sim20\\>$--$\\>$23, $B-I\\sim2\\>$--$\\>$3.5) in addition to these two components.  The location of this population of faint red objects in the color-magnitude diagram is as would be expected for red giant stars at the distance of M31.  The surface density of red giant candidates in the $R=19$~kpc M31 halo field is consistent with the findings of two recent {\\it Hubble Space Telescope\\/} studies.  The data indicate that M31's stellar halo is much denser and/or larger than that of the Galaxy: $(\\rho_{\\rm M31}^{\\rm  RGB}/\\rho_{\\rm MW}^{\\rm RGB})(\\Lambda/1.5)^{-\\nu}\\sim10$, where $\\Lambda$ is the ratio of the radial scale lengths of M31 and the Galaxy and $\\nu=-3.8$ is the assumed power law index of the density profile; in fact, M31's profile may be steeper than this ($\\nu<-3.8$).  The color and slope of the red giant branch in M31's outer halo are suggestive of a relatively metal-rich population, $\\rm[Fe/H]\\ga-1$, in agreement with the {\\it Hubble Space Telescope\\/} measurements. ",
        "introduction": "Studying the metallicity gradient of galactic spheroids is crucial for understanding their formation and evolutionary history.  The dissipational collapse model (Eggen et~al.\\ 1962; Larson 1974) predicts a strong metallicity gradient because of progressive chemical enrichment during the collapse of the protogalactic gas cloud.  The accretion model (Searle \\& Zinn 1978), on the other hand, predicts no strong gradient because star formation largely precedes assembly of the galactic spheroid.  The metallicity gradient of a halo can be measured in both the globular cluster system and in field stars; in fact, it is important to study both populations independently since they may be dynamically distinct from each other.  Globular clusters are easily observed out to large distances in the Galaxy because of their high luminosity and this allows a reliable determination of the metallicity gradient in the system of Galactic globular clusters.  Zinn (1993) finds evidence for two subsystems among clusters: an `old halo' population which displays a radial [Fe/H] gradient, a small age spread, and significant rotation, and a `young halo' population which displays no metallicity gradient, a large age spread, and very little rotation.  He associates the former population with the initial dissipational collapse of the Galactic halo (a la Eggen et~al.\\ 1962) and the latter with subsequent accretion of stellar systems (a la Searle \\& Zinn 1978).  On the other hand, Carney et~al.\\ (1990) find a mean metallicity of $\\rm[Fe/H]=-1.72$ and an absence of a strong metallicity gradient as a function of apogalacticon distance over the range $8~{\\rm kpc}\\leq{R}_{\\rm apo}\\leq40$~kpc in a proper motion selected survey of 227~Galactic field halo stars with velocities lagging the local standard of rest's circular velocity by at least 150~km~s$^{-1}$.  It seems logical to test the universality of the conclusions about galaxy formation which are being drawn from the single case of the Galaxy.  The Andromeda galaxy (M31) provides an external perspective of a large spiral similar to our own and yet is close enough for individual stars to be resolved.  Huchra et~al.\\ (1991) showed that there is evidence for a weak metallicity gradient in a sample of 150~M31 globular clusters, with a mean metallicity of $\\rm[Fe/H]=-1.2$, which is slightly higher than the mean value of $\\rm[Fe/H]=-1.4$ for Galactic globular clusters.  Over the last decade, several groups have tried to determine the metallicity of field red giant branch (RGB) stars in M31's spheroid, starting with the early work of Crotts (1986).  Mould \\& Kristian (1986) found a metallicity of $\\rm[Fe/H]=-0.6$ for stars in a field located at a projected distance of $R=7$~kpc along the minor axis, and assumed that the observed color spread was due to a large metallicity dispersion.  Mould (1986) failed to determine the metallicity at $R=20$~kpc, because the contaminating field galaxies could not be removed by the standard method of morphological selection and statistical subtraction. Pritchet \\& van den Bergh (1988) estimated $\\rm[Fe/H]=-1.0$ at $R=8.6$~kpc on the minor axis, and again attributed the large color spread to an intrinsic spread in [Fe/H].  Christian \\& Heasley's (1991) data in an $R=16$~kpc field around the globular cluster G219 suggested a large metallicity spread, with a mean value $\\rm[Fe/H]\\ga-1.0$, as did Davidge's (1993) study at $R=6.7$~kpc, Durrell et~al.'s (1994) study, and Couture et~al.'s (1995) study which targeted fields around five~M31 globular clusters.  While these studies have furthered our knowledge of M31's halo, sample contamination by distant field galaxies, foreground Galactic dwarfs, and M31 disk giants poses a serious obstacle in attempts to determine the metallicity of M31's outer halo ($R\\ga10$~kpc).   \\begin{figure} \\epsscale{0.8} \\plotfiddle{Reitzel.fig1.ps}{4in}{0}{70}{70}{-210}{-90} \\caption{Schematic plot showing the approximate location of the field studied in this paper, along with Holland et~al.'s (1996) G302 and G312 fields and Rich et~al.'s (1996) G1 field.  The size of each square symbol marking field position corresponds to the actual area of the field studied, but not to the field shape.  The ellipses indicate the orientation of M31's disk ($i=77^\\circ$, $\\rm PA=40^\\circ$, with finite thickness) and halo (5:3 flattening, same PA as disk); the sizes of the ellipses are arbitrary. The 10\\% change in the $\\rm cos(\\delta)$ projection factor from the top to the bottom of the plot has been ignored in this schematic representation. \\label{fig1}} \\end{figure}  Two recent studies of the M31 spheroid have taken advantage of the excellent angular resolution of {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) Wide Field/Planetary Camera~2 images to separate stars from distant field galaxies.  In the first study, Rich et~al.\\ (1996) analyze archival {\\it HST\\/} images of a field centered on the M31 globular cluster G1.  The line of sight to this field (Fig.~\\ref{fig1}) intersects M31's highly inclined disk ($i\\sim77^\\circ$; de~Vaucouleurs 1958) at a relatively large radius $R_{\\rm disk}\\approx{R}=40$~kpc on the major axis (the value of $R_{\\rm disk}$ could be different if M31's outer disk is warped).  The second study is that of Holland et~al.\\ (1996) who analyze deep {\\it HST\\/} images of two~fields around M31 globular clusters located 7.6~kpc (G302) and 10.8~kpc (G312) from the center roughly along the SE minor axis (Fig.~\\ref{fig1}).  Both studies find the mean metallicity of M31's halo RGB stars to be comparable to that of 47~Tuc ($\\rm[Fe/H]=-0.7$), with a spread of nearly 2~dex, suggesting that there was a greater degree of pre-enrichment during the assembly of M31's spheroid than in the case of the Galaxy's spheroid.  We will return to a more detailed discussion of these studies later in the paper (Sec.~4.2).  In this paper, we probe RGB stars in the M31 halo using an alternative method for discriminating between stars and faint galaxies; one that uses a combination of broadband $UBRI$ colors and morphological information (Gould et~al.\\ 1992, hereafter referred to as GGRF).  Early results have been reported by Reitzel et~al.\\ (1996a,b).  The field of study presented here is at a projected distance of $R=19$~kpc from M31's center along the minor axis (only Rich et~al.'s G1 field is at a comparable radial distance out in the spheroid assuming it has a 5:3 flattening) and the line of sight to the field intersects M31's disk at $R_{\\rm disk}\\approx75$~kpc, further out than in any of the previous studies (Fig.~\\ref{fig1}).  The area of the field in this study, 210~arcmin$^2$, is more than 60~times larger than each of the Holland et~al.\\ fields, and more than 40~times larger than the Rich et~al.\\ G1 field.  Moreover, the field stars targeted in this study are well away from any of M31's globular clusters.  The observations and data reduction procedure are described in Sec.~2.  The technique  used to isolate M31 halo RGB stars from background field galaxies and foreground Galactic dwarfs is described in Sec.~3.  In Sec.~4, we discuss the properties of M31's spheroid in the context of the recent {\\it HST\\/} studies and outline plans for future work.  The main points of the paper are summarized in Sec.~5. ",
        "conclusions": "\\subsection{Estimating the Metallicity of M31's Outer Spheroid}  The metallicity of the population of M31 halo stars detected is estimated by comparing the color-magnitude distribution of the $UBRI$- and FWHM-selected, statistically-subtracted sample of objects to the RGB fiducials for the Galactic globular clusters M5 and 47~Tuc in Fig.~\\ref{fig8}(f). Clearly the fiducial for M5 is not a good match as most of the M31 halo field objects brighter than $I=23$ are redder, and hence more metal rich, than M5 ($\\rm [Fe/H]=-1.4$).  The clump of faint ($I>23$), blue ($B-I\\sim1\\>$--$\\>$2) objects in this plot is produced at least in part by residual contamination by distant faint blue field galaxies [compare Figs.~\\ref{fig8}(c) and (f)].  The fiducial for 47~Tuc is a better match to the data, and yet there is a substantial spread about this fiducial.  A substantial fraction of the objects with $I<23$ lie to the right of the 47~Tuc fiducial which indicates that the mean metallicity of M31's outer halo is comparable to, and possibly a bit higher than, that of 47~Tuc ($\\rm [Fe/H]=-0.70$), with a fairly large spread in color.  The mean photometric errors estimated from Poisson noise are $\\sigma_{B-I}\\simeq0.02$ for objects near the tip of the M31 RGB ($I\\sim20.5$) and $\\sigma_{B-I}\\simeq0.18$ near the limit of the sample ($I\\sim24$).  The observed spread in color is too large to be accounted for simply in terms of photometric error.  Despite our efforts to remove all possible contaminants, some of the observed spread could be due to residual contamination from foreground Galactic stars and background faint compact galaxies.  The IASG model has not been thoroughly tested at faint magnitudes; moreover, Poisson fluctuations in the field galaxy counts cause differences between the comparison and M31 halo fields.  These possible contaminants cannot be entirely responsible for the observed color spread however as there is no reason why the contaminants should contribute primarily near the observed CMD location of the M31 RGB.  In fact, compact galaxies contribute primarily at faint magnitudes ($I\\ga22$) and blue colors ($B-I\\leq2$) as can be seen by comparing the $UBRI$- and FWHM-selected comparison field sample in Fig.~\\ref{fig8}(b) with the IASG model prediction for this field in Fig.~\\ref{fig8}(d).  The uncertainty of the IASG model at faint magnitudes is harder to quantify; some Galactic stars may still be contaminating even the foreground subtracted CMD in Fig.~\\ref{fig8}(f).  We conclude that much of the spread observed in the RGB concentration is due to a true spread in the metallicity of M31 halo stars.  Future followup spectroscopy will test this conclusion (Sec.~4.3).  \\subsection{Comparison of this Study to Recent HST Studies}  In this section, we compare the final, statistically-subtracted sample of RGB candidates in the M31 halo CMD to those found in the {\\it HST\\/} studies of Holland et~al.\\ (1996) and Rich et~al.\\ (1996).  Fig.~\\ref{fig1} shows the locations of all the fields in relation to M31's disk and halo (schematic representation).  We adopt the authors' plausible hypothesis that the vast majority of objects seen in Rich et~al.'s $R=40$~kpc (G1) CMD and in Holland et~al.'s $R=7.6$~kpc (G302) and 10.8~kpc (G312) CMDs are M31 field halo stars (see details below).  We use the External Galaxy Model (EGM) developed by Hodder (1995) to estimate the number of halo and disk stars expected to be seen in each of these fields.  The EGM is an adaptation of the Bahcall \\& Soneira (1984) Galaxy model designed to give star counts and color distributions for any field in an external spiral galaxy using a de~Vaucouleurs profile for the halo and an exponential disk. \\begin{table} \\tablenum{2} \\begin{tabular}{lrrrrrrr} \\multicolumn{8}{c}{\\smcap Table~2.} \\\\ \\multicolumn{8}{c}{Observed Star Counts and EGM Predictions} \\\\ \\hline \\hline  \\multicolumn{1}{c} {Field} & \\multicolumn{1}{c} {Area} & \\multicolumn{1}{c} {$X^a$} & \\multicolumn{1}{c} {$Y^a$} & \\multicolumn{2}{c} {Observed} & \\multicolumn{2}{c} {EGM Predictions$^b$} \\\\   \\multicolumn{1}{c} {} & \\multicolumn{1}{c} {[arcmin$^2$]} & \\multicolumn{1}{c} {[kpc]} & \\multicolumn{1}{c} {[kpc]} & \\multicolumn{1}{c} {$I$=20--22} & \\multicolumn{1}{c} {$I$=22--24} & \\multicolumn{1}{c} {$I$=20--22} & \\multicolumn{1}{c} {$I$=22--24} \\\\  \\hline  This Study$\\,\\>$ \\hfill & 210.0 & \\hfill $-1.7$ & \\hfill $-18.7$ & \\hfill 122$^c $ & \\hfill 543$^c$ & \\hfill 421 (421) & \\hfill 6141 (6140) \\\\ G302$^d$ \\hfill & \\hfill 3.3 & \\hfill $-2.5$ & \\hfill $-6.7$ & \\hfill 339$\\,\\>$ & \\hfill 1756~  & \\hfill 132$\\,\\>$(119) & \\hfill 1955 (1748)\\\\ G312$^d$ \\hfill & \\hfill 3.3 & \\hfill $+0.5$ & \\hfill $-11.1$ & \\hfill 82~$\\,$ & \\hfill 436$\\,\\,$ & \\hfill 33$\\,\\>$(32)$~\\,$ & \\hfill 486 (479)$~\\,$ \\\\ G1$^e$ \\hfill & \\hfill 1.1 & \\hfill $+33.6$  & \\hfill $+4.6$ & \\hfill 33~$\\,$ & \\hfill 140$\\,\\,$ & \\hfill 10$\\,\\,$(6)~~~ & \\hfill 157$~$(91)~~~ \\\\  \\hline  \\end{tabular}  \\begin{itemize} \\item[$^a$]{($X$, $Y$)~=~Projected distance of field center from M31's center along its major and minor axes, respectively} \\item[$^b$]{Numbers in parentheses indicate spheroid stars} \\item[$^c$]{The star selection efficiency for our $R=19$~kpc sample is estimated to be about 30\\% for the brighter apparent magnitude range and probably as low as 10\\%$\\>$--$\\>$20\\% for fainter stars} \\item[$^d$]{From the study by Holland et~al.\\ (1996)} \\item[$^e$]{From the study by Rich et~al.\\ (1996)}  \\end{itemize} \\end{table}  There is a large number of free parameters in the EGM (disk and halo scale lengths, density ratio, halo axis ratio, and overall normalization), but a plausible model based on 47~Tuc's red giant luminosity function and color distribution, a disk exponential scale length of 5.5~kpc, a halo effective radius of 2.6~kpc, and an axis ratio for the halo of~0.6 can be normalized to match the star counts observed in the {\\it HST\\/} fields to within $\\sim$3\\%. Furthermore, the fraction of M31 disk stars predicted by the EGM for the two Holland et~al.\\ fields agrees with the upper limits to the fractional disk contamination estimated by the authors: 0.1 for G302 and 0.03 for G312.  The halo effective radius of 2.6~kpc for the major axis corresponds to a minor axis effective radius of 1.6~kpc for the flattened halo.  This is consistent with the minor axis effective radius of 1.3~kpc found by Pritchet \\& van~den~Bergh (1994).  The EGM predicts about 6500~RGB stars with $20<I<24$ within the 210~arcmin$^2$ effective area of the KPNO image of the M31 halo. Most of these are expected to be halo giants with negligible contamination from M31's disk (0.01\\%).   Table~2 presents the observed number of stars and EGM predictions for this study and for the three {\\it HST\\/} fields.  The surface density of M31 stars predicted by the EGM in the $R=19$~kpc field we have studied is lower than in any of the three {\\it HST\\/} fields, even Rich et~al.'s G1 field at $R=40$~kpc.  A 5:3 flattened halo contributes about equally to the G1 field and ours, but M31's disk contributes an additional 70\\% (40\\% of total) in the G1 field.  Contrary to the impression one may get from Fig.~\\ref{fig1}, the degree of disk contamination is much higher in G1 than in the inner Holland et~al.\\ (G302) field.  Even though the G302 field is expected to contain {\\it more\\/} disk stars than G1, this is more than compensated by the much larger number of halo stars in G302, making the fractional disk contamination in G302 quite small.  This is simply a reflection of the fact that galaxy halos have steeper density profiles than the typical disk exponential law.  The M31 disk RGB fraction in G1 could be even higher if the halo is less flattened ($b/a>0.6$) and/or if a thick disk is present.  Thus, the large spread in the M31 halo metallicity inferred by Rich et~al.\\ may in fact be due primarily to disk contamination (possibly including stars with super-solar [Fe/H]).  After FWHM and $UBRI$ selection and statistical subtraction, we detect 665~objects down to a limit of $I<24$ in the M31 halo field [see Fig.~\\ref{fig8}(f)], whereas the EGM predicts about 6500.  This would seem to indicate a $\\sim10\\%$ selection efficiency, but the efficiency is higher than this over most of the magnitude range.  About $\\sim30$\\% of the number predicted by the EGM is observed in the range $20<I<22$, which is close to estimated efficiency of the star selection process at relatively bright magnitudes (as discussed in Sec.~3.6).  At fainter magnitudes the selection efficiency drops for a variety of reasons: non-gaussian photometric errors (effect of crowding on isophotal definition, sky measurement errors, etc.) and increased FWHM measurement error (causing stars to scatter beyond $\\theta_{\\rm crit}$).  Overall, the number of objects seen in our field is roughly consistent with the EGM model prediction given the selection efficiency.  The estimated metallicity of M31's outer halo from our study (Sec.~4.1) is consistent with that found in the recent {\\it HST\\/} studies. Holland et~al.\\ (1996) find a spread in metallicity of $\\rm-2\\la[Fe/H]\\la-0.2$ with the majority of stars having [Fe/H]$\\simeq-0.6$. The contamination due to M31 disk stars is estimated using the EGM to be $\\sim0.1$ for the G302 field and $\\sim0.03$ for the G312 field.  This, along with the estimated foreground Galactic contamination of only 15~disk stars, cannot completely explain the population of very red (high metallicity) objects observed in the G302 and G312 fields.  If the disk of M31 is warped, however, the fractional disk contamination could be higher in these Holland et~al.\\ fields.  Rich et~al.\\ (1996) find the field population around G1 to be as metal rich as 47~Tuc, but warn that the apparent high metallicity could also be due to the sample containing objects of intermediate age.  As discussed above, contamination by M31's disk is expected to be significant in the G1 field.  \\subsection{Density and Size of the M31 Halo}  The observed RGB counts in the Holland et~al.\\ (1996) {\\it HST\\/} minor axis fields and in our $R=19$~kpc minor axis field can be used to constrain the physical parameters of M31's halo.  The halo is assumed to be oblate and viewed edge on---i.e.,~with its symmetry axis in the plane of the sky.  A (deprojected) power law index of $\\nu=-3.8$ fits the surface densities observed in the $R=7.6$~kpc (G302) and $10.8$~kpc (G312) Holland et~al.\\ fields derived from stars with $20<I<23.5$.  This slope is a bit shallower than but consistent with the range of $\\nu=-4$ to $-5$ found by Pritchet \\& van~den~Bergh (1994) over the (minor axis) radial range of $R=3\\>$--$\\>$20~kpc based on stars within 2~mag of the tip of the RGB. Pritchet \\& van~den~Bergh find that M31's halo density profile steepens with increasing radius (in a log-log plot), so it is not surprising that a single power law fit to the inner halo ($R<10.8$~kpc) is shallower than one that includes the outer halo.  For a power law density profile, the column density at a projected radial distance $R$ along the minor axis is given by:  \\begin{equation} N=\\int_{-\\infty}^{\\infty}n_0\\left[\\frac{x^{2}+(R/c)^{2}}{R_0^{2}}\\right]^{\\nu/2}dx \\end{equation} \\begin{equation} N=n_0\\left(\\frac{R}{c}\\right)\\left(\\frac{R_0\\,c}{R}\\right)^{-\\nu}G(\\nu) \\end{equation}  \\noindent where $c=0.6$ is the halo axial ratio, $x$ is the distance along the line of sight, and  \\begin{equation} G(\\nu)\\equiv\\frac{\\Gamma(-\\nu/2\\>-\\>1/2)\\,\\Gamma(1/2)}{\\Gamma(-\\nu/2)} ~~~. \\end{equation}  \\noindent The volume density $n_0$, at a radial distance $R_0$ in the equatorial plane of the halo, may thus be inferred from the observed surface density:  \\begin{equation} n_0={{\\frac{N\\,c}{R\\,G(\\nu)} \\left(\\frac{R}{R_0\\,c}\\right)^{-\\nu}}} \\end{equation}  The halo volume density is calculated in the apparent magnitude range $20<I\\leq23.5$ to avoid red clump stars and to directly compare to Morrison's (1993) measurement of the local Galactic halo stellar density.  A projected density of $N=7870$~kpc$^{-2}$ is derived from counts of stars (1290) in this $I$ magnitude range in the inner Holland et~al.\\ field (G302).  The area of this field (two~Wide Field Camera CCDs) is 3.3~arcmin$^2$ which corresponds to 0.17~kpc$^2$ at the distance of M31.  The inner (7.6~kpc) Holland et~al.\\ field, G302, is not exactly on the minor axis; the ellipse with $c=0.6$ that intersects this field has a semi-minor axis length of $R=6.8$~kpc, and this is the value used in the above equations.  As M31 is about $\\Lambda\\approx1.5$~times larger than the Galaxy, we choose $R_0=12$~kpc in order to calculate the halo density at a position in M31 comparable to the solar location in the Galaxy.  The derived M31 halo density at $R_0=12$~kpc is $n_0=341$~kpc$^{-3}$.  Thus, $n_0$ is much larger than the local Galactic halo density of 36~kpc$^{-3}$ for stars with $M_V\\leq0.5$ (Morrison 1993).  This indicates that M31's halo is much denser, and/or larger than that of the Galaxy, $(\\rho_{\\rm M31}^{\\rm RGB}/\\rho_{\\rm MW}^{\\rm RGB})(\\Lambda/1.5)^{-\\nu}\\sim9.5\\pm0.3$, where $\\Lambda$ is the ratio of the characteristic radial scale lengths of M31 and the Galaxy, $\\nu\\sim-3.8$ is the power law density profile slope, and the (formal) error bar is based on Poisson error in the star counts in the inner Holland et~al.\\ field, G302.  The uncertainty in the overdensity estimate is substantially higher than the formal error bar quoted above when one takes into account the possible range of profile slopes, $\\nu$ (see below).  The actual density of M31's inner halo is likely to be even higher than this estimate:  We have adopted a single power law profile in the above integral, with a slope of $\\nu=-3.8$ based on M31's inner halo, but the outer profile is known to steepen both from the counts in our $R=19$~kpc field and from Pritchet \\& van~den~Bergh (1994); this leads to an overestimate of the column density for a given spatial density $n_0$, or to an underestimate of $n_0$ since we normalize to the {\\it observed\\/} column density.  At any rate, it is clear than the halo of M31 is significantly denser/larger than the halo of the Galaxy.  The density profile of the outer halo of M31 (beyond a minor axis radius of 10~kpc) has a slope of $\\nu\\approx-5$, as indicated by the EGM de~Vaucouleurs spheroid fit to the observed counts at $R=10.8$~kpc (G312, Holland et~al.\\ 1996) and $R=19$~kpc (this study), and by Pritchet \\& van~den~Bergh's (1994) study.  This is steeper than the profile of outer Galactic halo which may be approximated by a power law index in the range $\\nu_{\\rm MW}=-3.0$ to $-3.5$ (Zinn 1985; Preston et~al.\\ 1991; Gould et~al.\\ 1998).  There are significant uncertainties though in these measurements of the profile of the Galaxy's outer stellar halo, and the halo of M31 is not well fit by a power law but rather by a de~Vaucouleurs law (Pritchet \\& van~den~Bergh 1994), so directly comparing the power law indices may be misleading.  \\subsection{Future Work/Followup Spectroscopy}  Although it is becoming increasingly clear that the halo of M31 is fairly metal rich, an accurate mean metallicity and spread remain to be determined. Contamination from foreground Galactic stars and background field galaxies, as well as from disk RGB stars in M31, is still non-negligible.  The combined color- and morphological-selection technique described in this paper can be used effectively to study M31's inner halo where the fractional contamination is expected to be lower than in our $R=19$~kpc field.  High angular resolution imaging (e.g.,~with {\\it HST\\/}) can help reduce background galaxy contamination.  We are currently undertaking followup Keck/LRIS spectroscopy of the brightest decade of M31 halo RGB candidates in the $UBRI$- and FWHM-selected sample [Fig.~\\ref{fig6}(d)] (Guhathakurta \\& Reitzel 1998).  Candidates for spectroscopic observation are selected without regard to color, so as not to introduce any metallicity bias.  We do not select objects from the statistically-subtracted version of the CMD as it is not complete. Spectroscopy should allow kinematic confirmation of the identity of M31 RGB stars: galaxies are eliminated altogether on the basis of their redshift and M31 stars are quite well separated from foreground Galactic disk stars ($v^{\\rm sys}_{\\rm M31}\\sim-300$~km~s$^{-1}$ vs $v_{\\rm MW~disk}\\approx0$~km~s$^{-1}$).  A broad versus narrow distribution of radial velocities will distinguish M31 halo stars from disk stars on M31's minor axis, respectively.  In addition to removal of residual contaminants, spectroscopy allows a direct measurement of the metallicity of each star using the near infrared \\Catwo\\ line strengths, W$_{\\rm Ca}$.  A well developed empirical calibration method exists to determine [Fe/H] from W$_{\\rm Ca}$ and the brightness of the star $\\delta{I}$ relative to the horizontal branch level (Olszewski et~al.\\ 1991; Armandroff \\& Da Costa 1991; Suntzeff \\& Kraft 1996)."
    },
    "9804/astro-ph9804035_arXiv.txt": {
        "abstract": "I present a matched-wave asymptotic analysis of the driving of solar oscillations by a general localised source. The analysis provides a simple mathematical description of the asymmetric peaks in the power spectrum in terms of the relative locations of eigenmodes and troughs in the spectral response. It is suggested that the difference in measured phase function between the modes and the troughs in the spectrum will provide a key diagnostic of the source of the oscillations. I also suggest a form for the asymmetric line profiles to be used in the fitting of solar power spectra.  Finally I present a comparison between the numerical and asymptotic descriptions of the oscillations. The numerical results bear out the qualitative features suggested by the asymptotic analysis but suggest that numerical calculations of the locations of the troughs will be necessary for a quantitative comparison with the observations. ",
        "introduction": "Helioseismology is primarily concerned with the peaks in the solar oscillation spectrum which are identified with acoustic eigenmodes of the Sun. However the whole spectrum, including the variation between the mode peaks, contains useful information about the Sun's structure and dynamics. Moreover, the asymmetries in the \\pmode~lines which were first identified by \\markcite{D93} Duvall et al. (1993) can bias the fitting of profiles to the peaks and hence introduce errors in the determination of the frequencies. The most recent space-based spectra from SOHO-MDI \\markcite{Netal98} (Nigam et al. 1998) have confirmed the existence of a strong asymmetry. Thus the study of \\pmode~line asymmetries is prompted both by the need to improve on traditional data analysis techniques and by the hope of obtaining new scientific information from them.  The occurrence of highly asymmetric line profiles is a common feature of models in which the oscillations are driven by a localised source \\markcite{G92,G93,G95}(Gabriel 1992, 1993, 1995). \\markcite{AK96}Abrams and Kumar (1996) explored the nature of the asymmetry in a simple model using analytical techniques and then applied these ideas to the interpretation of numerical calculations of the line asymmetry in a real solar model, including the variation of the asymmetry with frequency and mode degree and estimates of the error introduced into the determination of the frequencies by the neglect of the line asymmetry.  \\markcite{RB97}Rast \\& Bogdan (1997) emphasised the sensitivity of the asymmetry to the depth and type of the source. In particular they introduced the approach that the observed asymmetry can best be approached through understanding the relative locations of the peaks and troughs in the spectrum, where the location of the troughs can be determined (in some circumstances) by solving an eigenproblem for a solar model truncated from above at the source depth. Although the calculations presented here will result in an estimate for the whole spectrum, I will follow the example of Rast \\& Bogdan in concentrating particularly on the relative locations of the peaks and troughs as a key diagnostic indicator of the source of excitation.  \\markcite{RV95}Roxburgh \\& Vorontsov (1995) used a form of asymptotic analysis to model the power spectrum. The analysis presented here differs from theirs in that it uses a matched wave technique in which the matching across the turning point(s) is treated explicitly. Moreover, I consider a general form of the source term and obtain an explicit expression for the asymptotic line profile. It is found that the power spectrum can be represented in a simple form even for a general source, and this should guide the analysis and interpretation of the peaks and troughs in the solar \\pmode~spectrum.  I begin in \\S\\ref{illus} with a very simple illustrative model which nevertheless captures some important aspects of the dynamics. In \\S\\ref{strat} I show how one may carry out a matched-wave asymptotic analysis for a stratified layer with a source containing both monopole and dipole terms. I am predominantly interested in the lowest frequency modes, for which the asymmetry is most easily observed. However, asymptotic results for intermediate and high-frequency oscillations are also presented. In \\S\\ref{results} I evaluate the resulting asymptotic expressions for a simple model and compare them with a direct numerical calculation. ",
        "conclusions": "Typically, helioseismology has concentrated on analysis of the eigenmode frequencies of the Sun and relatively little attention has been paid to the rest of the spectrum. With the advent of the SOI/MDI instrument on SOHO and the GONG network, noise levels in measured spectra are now sufficiently low that it is not only possible to make useful measurements between the mode peaks but {\\it necessary}\\/ to do so since it is increasingly clear that the line asymmetry can produce systematic errors in the determination of {\\it p}\\/-mode frequencies if it is not accounted for properly.  A puzzle in the data has been the observation originally due to Duvall et al (1993), and now apparently confirmed from SOI-MDI (Nigam et al. 1998), that the sign of the measured asymmetry is oppositely directed in doppler and intensity measurements. The doppler measurements imply a dominant monopole source (a negative asymmetry in the Abrams \\& Kumar terminology) while the intensity measurements would suggest a dipole source. Rast \\& Bogdan (1997) show that, regardless of the form of the source function, one would expect to see the same sign of asymmetry with observations made in different variables, since the locations of the peaks and troughs are the same for all variables. Nigam et al. have shown that the phenomenon can be explained if the measured signal in either doppler or intensity is contaminated by a noise source which is correlated with the source, for example if the measured signal contains a contribution coming directly from the source itself. The question is then raised as to which of the doppler measurement and the intensity measurement actually constitutes a better measurement of the true acoustic spectrum. One might argue that the doppler signal is probably less contaminated since it has a lower overall background level, but it is certainly possible that both the intensity and doppler measurements contain some leakage from the source, with an amplitude and phase which are not yet precisely known. Evidently the possibility that the location of the spectral minima is shifted due to such leakage must complicate the interpretation of the line-asymmetry in terms of source location and type, and further work will be required to determine whether the combined doppler and intensity measurements can be used to determine the true acoustic spectrum in a model-independent fashion.  As originally pointed out by Rast \\& Bogdan, the asymmetry is best approached by consideration of the minima in the spectrum, whose location is determined by the depth and type of the source of the oscillations. In this work I have determined an asymptotic expression for the frequency spectrum, from which it is possible to determine the locations of the minima. The sign of the line asymmetry is determined by whether the nearest minimum lies above or below the eigenfrequency. The former case corresponds (at low frequency) to a pure monopole source and the latter to a pure dipole. The amount of asymmetry depends of the location of the source. The higher the source is placed above the upper turning point of the oscillation the closer the trough will lie to the eigenmode and therefore the greater the asymmetry. Thus the extent of the asymmetry should be a useful diagnostic of the properties of the turbulent superadiabatic layer in which the source of the solar \\pmodes~is believed to lie.  Numerical results confirm the qualitative properties of the asymptotic results. Also for the case of a pure monopole or dipole, or where the two components of the source are in phase, the exact locations of the nulls (including those above the acoustic cutoff frequency) can be determined by solving numerically the oscillation equations subject to the boundary condition \\begin{equation} {\\dd\\Psi\\over \\dd z}- {a_0\\over a_1} |K| \\Psi=0 \\label{bc1} \\end{equation} at $z=z_0$. This should be relatively straightforward to implement in most existing oscillation codes. One can thus determine numerically the phase-function difference $\\alpha_t-\\alpha_m$ which is predicted by the asymptotic analysis to be a key diagnostic of the structure and location of the source.  Since the nulls are independent of the superficial structure the asymmetry could, in principle, provide a tool for differentiating between superficial and more-deeply seated changes in the Sun. In particular, solar-cycle variations occurring above the source location, in the lower atmosphere for example, would change the mode frequencies but not the trough frequencies and could thus be visible as changes in the asymmetry, while changes below the source would change the frequency of both the mode and the trough. In practice such changes are likely to be extremely difficult to observe, not least because the asymmetry is most evident for lower frequency modes for which the solar-cycle dependent changes are small.  We have thus reached the stage where it is possible to envisage a program for the re-analysis of helioseismic data to make full use of the information contained in the \\pmode~line profiles. In particular it would be highly desirable to re-analyse the spectral data using parameterised line-profiles such as equation (\\ref{profile1}) or equation (\\ref{profile2}) in order to obtain improved estimates of the eigenfrequencies simultaneously with estimates for the spacing between the maxima and minima. From these we could obtain an observational $\\alpha_t-\\alpha_m$ curve which could be calibrated against theory using nulls calculated from an eigenvalue calculation with boundary condition (\\ref{bc1}). Ultimately one would hope that the results of such a calibration could be brought into line with predictions made from numerical simulations of the solar surface layers ({\\it e.g.}\\/\\markcite{SN89,SN91,Netal96} Stein \\& Nordlund 1989, Stein \\& Nordlund 1991, Nordlund, Stein \\& Brandenburg 1996) to provide a solidly grounded understanding of this highly complex region of the Sun."
    },
    "9804/astro-ph9804173_arXiv.txt": {
        "abstract": "The properties of X-ray Nova XTE J1755--324, observed with RXTE observatory in 1997 are reported. The lightcurve of the source was typical for X-ray Novae, except for somewhat shorter decay time scale and the time elapsed between the primary and subsequent secondary and tertiary outbursts. At the peak of the lightcurve the source had two-component spectrum, typical for the bright X-ray Novae, with a characteristic disk blackbody temperature $T_{in}\\sim0.8$ keV and a photon index of the hard power law tail $\\alpha\\sim$ 2.0. The peak luminosity can be roughly estimated as $L\\sim10^{38}$ erg/s (0.1--100 keV, assuming 8.5 kpc distance). A notable and peculiar feature of the spectral evolution of the source was a short, $\\sim 10$ days long, episode of increased hardness, occurred shortly before the tertiary maximum of the light curve. During the last RXTE pointed observations $\\sim 100$ days after the primary outburst the source  has been found in the hard spectral state with luminosity $\\sim$few$\\times 10^{36}$ erg/s. The overall pattern of the temporal and spectral evolution of XTE J1755--324 resembles in general that of ``canonical'' X--ray Novae (e.g. Nova Muscae 1991) and suggests that the compact object in the binary system is a black hole. ",
        "introduction": "XTE J1755--324 was discovered by All--Sky Monitor (ASM) aboard Rossi X-ray Timing Explorer (XTE) on July 25, 1997 and localized by Proportional Counter Array (PCA) at the position with coordinates R.A.=17$^h$55$^m$28$^s$.6, Dec.=--32\\fdg28\\arcmin39\\arcsec (J2000). The light curve of the source was typical for X-ray Novae (e.g. X-ray Nova Muscae 91, X-ray Nova Vulpeculae 88, see \\cite{tanaka} for a review) with a quick rise within a few days and an exponential decay with an e--folding time of $\\sim$30--40 days. The spectrum of the source at the peak of the lightcurve was complex, with a soft component having approximately multicolor blackbody disk spectrum with $T_{in}\\sim0.8$ keV and a hard power law tail extending above 10 keV. Preliminary analysis of the data of BeppoSAX observations of the Galactic Center region in Sep. 1997 has not revealed X-ray bursts from the source (Ubertini 1997, private communication). Radio observations of the region of the sky containing the source performed on Aug 18, 1997 have not found significant radio emission from XTE J1755--324 down to $\\sim0.3$ mJy (\\cite{radio}). The optical extinction at the direction of the source corresponding to $N_H\\sim3.7\\times10^{21}$ cm$^{-2}$ (\\cite{nh}) is $A_v\\sim$2, however to our knowledge no observations of the source at optical wavelengths have been reported. ",
        "conclusions": "The pattern of spectral and temporal evolutions of XTE J1755--324 bears several features in common with other X--ray Novae, e.g. X-ray Nova Vulpeculae 88 (\\cite{mir-kvant}), Nova Muscae 1991 (\\cite{miyamoto}, \\cite{tim}), KS/GRS 1730--312 (\\cite{tsp}), GRS 1739--278 (\\cite{kbor1739}): \\begin{enumerate} \\item The shape of the X--ray light curve (short risetime, nearly exponential decay, secondary and tertiary outbursts -- Fig.\\ref{asmcurve}) is frequently observed for X--ray Novae (\\cite{tanaka}). \\item The two component spectrum and the low level of aperiodic variability of the X--ray flux observed at the peak of the light curve are well established  features of the very high state of the black hole candidates and are often observed at the peak of the light curve of other X--ray Novae.  \\indent The spectrum of XTE J1755--324 was dominated by the soft component which contributed more than 90\\% to the source luminosity. The shape of the soft component suggests that it may originates from the optically thick part of the accretion disk. Spectral fit by a simple multicolor blackbody disk model indicates that  the inner radius of the optically thick disk was sufficiently close to the compact object. However, one should bear in mind that the value of the inner disk radius derived above depends on the spectral model, the binary system inclination angle and assumed source distance.  Based on the source spectrum approximated by the two-component model consisting of a multicolor blackbody disk model and a hard power law component, the 3--30 keV  luminosity near the peak of the light curve is $L_{3-30 keV}\\sim1.3\\times10^{37}$ erg/s for 8.5 kpc distance. The total absorption corrected luminosity of the source in the 0.1--25.0 keV energy band can be estimated as  $L_{0.1-25 keV}\\sim1\\times10^{38}$ erg/s. This value is close to the luminosity of other black hole candidates in the very high state, which might indicate that the assumption about the source distance is roughly correct.  \\item Anticorrelated behavior of the X--ray flux and the 5.0--12.2 keV to 1.3-3.0 keV  hardness ratio before and during some time after the primary outburst (Fig.\\ref{asmcurve}, the bottom panel).  The hardness ratio plot shows that the spectrum of XTE J1755--324 softened during the initial rise of the flux. Similar behavior was observed for other X-ray Novae -- GRS/GS 1124--64 (Nova Muscae 1991, \\cite{miyamoto}), KS/GRS 1730--312 (\\cite{tsp}), GRS 1739--278 (\\cite{kbor1739}), and is likely due to the increase of the relative contribution and the absolute luminosity of the soft spectral component caused by increase of the mass accretion rate and corresponding increase of the disk temperature.  After the peak of the light curve  the hardness ratio increased gradually and fairly smoothly with a brief  excursion around TJD 10708. The gradual increase of the hardness ratio after the primary outburst is also seen in the spectral evolution of other X--ray Novae. Although no spectrally resolved data is available in the case of XTE J1755--324, we may suppose that it is related to decrease of the mass accretion rate and corresponding decrease of the disk temperature (that is usually followed by a transition to the low spectral state).   \\item Transition to the low spectral state during decay.  The transition of the source to the low spectral state itself was not traced by the pointed instruments aboard RXTE, but during November 1997 observations the source has been found in the low spectral state according to both spectral (Fig.\\ref{spectrum}, the right panel) and aperiodic variability (Fig.\\ref{pds}) properties. The 4-100 keV luminosity was $L_{4-100 keV}\\sim5\\times10^{36}$ erg/s (Nov. 12, 1997,  PCA and HEXTE data, the HEXTE normalization was adjusted to that of PCA).   \\end{enumerate}     \\medskip     A notable and apparently peculiar feature of the source evolution was a brief excursion seen on the hardness ratio plot (Fig.\\ref{asmcurve}) around TJD 10708. Comparison of July-Aug. 1997 and Sep. 1997 spectra of XTE J1755--324 observed with RXTE/HEXTE and GRANAT/SIGMA (Paul et. al 1997) in the high energy domain further confirms a fact of strong spectral evolution  (see Fig.\\ref{spectrum}). The RXTE/HEXTE observation near the peak of the light curve on Aug. 1, 1997 gave the hard X--ray flux $F_{40-75 keV}\\sim9\\times10^{-11}$ ergs/s/cm$^2$. The relative contribution of the hard power law component to the total flux in the 3--30 keV energy band was at the level of $\\sim$10\\%. The hard X-ray flux from the source measured by GRANAT/SIGMA near the peak of the hardness is $F_{40-75 keV}\\approx5\\times10^{-10}$ ergs/s/cm$^2$ (\\cite{sigmaiauc}), i.e. the 40--75 keV flux from the source increased by a factor of $\\sim$5. At the same time according to RXTE/ASM  data  the 1--12 keV flux decreased by a factor of $\\sim$4--5 (Fig.\\ref{asmcurve}). Obviously the observed spectral changes can not be accounted for in terms of a simple scaling of the soft spectral component  according to the luminosity change (see p.3 above). Moreover, comparison of the ASM spectra obtained before/after and during this event (Fig.\\ref{spectrum}, the middle panel) shows clearly, that not only luminosity in the high energy band increased but also the luminosity below $\\sim 5$ keV decreased by a factor of $\\sim 2-3$. No evidence of the soft spectral component of the type observed just before and immediately after the event is seen in the ASM data taken at the peak of hardness.   Assuming a power law spectrum with a photon index $\\alpha=2.0$ and a high energy cutoff (Comptonization spectrum with kT$\\sim20$ keV and $\\tau \\sim4$) which roughly describes ASM and SIGMA data in the 1-150 keV range, (solid line in the middle panel in Fig.\\ref{spectrum}) the 1.3--150 keV luminosity is $\\sim2\\times10^{37}$ erg/s, which is close to typical low spectral state luminosity of Cyg X-1 and 1E1740.7--2942 (e.g. Sunyaev \\& Truemper 1979, Sunyaev et. al 1991).  In general, the shape of the spectrum, observed with RXTE/ASM and GRANAT/SIGMA in the middle of September 1997 is typical for the low state of the black hole binaries except for the noticeably lower value of the Comptonization temperature. A possible explanation of the event might be that it is due to a spectral transition to the hard state followed by a transition back to the soft state caused by an increase of the mass accretion rate corresponding to the tertiary maximum of the light curve.  Such an event can ocure when the accretion rate is near a critical level: e.g. a hard-soft-hard transition in Cyg X--1 (Zhang et al., 1997). Crucial data for understanding the nature of the event might be the timing information which unfortunately is absent -- no pointed observations were performed by RXTE at that time.  The pattern of the spectral and temporal evolution of XTE J1755--324 is similar to that of ``canonical'' X--ray Novae  of dynamically proven black hole candidates and suggests that the binary system XTE J1755--324 also harbors a black hole."
    },
    "9804/astro-ph9804021_arXiv.txt": {
        "abstract": "We explore mechanisms for the regulation of star formation in dwarf galaxies. We concentrate primarily on a sample in the Virgo cluster, which has HI and blue total photometry, for which we collected H$\\alpha$ data at the Wise Observatory. We find that dwarf galaxies do not show the tight correlation of the surface brightness of H$\\alpha$ (a star formation indicator) with the HI  surface density, or with the ratio of this density to a dynamical timescale, as found for large disk or starburst galaxies. On the other hand, we find the strongest correlation to be with the average blue surface brightness, indicating the presence of a mechanism regulating the star formation by the older (up to 1 Gyr) stellar population if present, or by the stellar population already formed in the present burst. ",
        "introduction": "The star formation (SF) is a fundamental process in the evolution of galaxies and is far from being well understood. The SF is usually characterized by the initial mass function (IMF) and the total SF rate (SFR), which depends on many factors such as the density of the interstellar gas, its morphology, its metallicity, {\\it etc.} According to Larson (1986), four major factors drive star formation in galaxies: large scale gravitational instabilities, cloud compression by density waves, compression in a rotating galactic disk due to shear forces, and random cloud collisions. In galaxies with previous stellar generations additional SF triggers exist, such as shock waves from stellar winds and supernova explosions. In dense environments, such as clusters of galaxies and compact groups, tidal interactions, collisions with other galaxies, ISM stripping, and cooling flow accretion probably play some role in triggering the SF process. The triggering mechanisms were reviewed recently by Elmegreen (1998).  While ``global'' phenomena, such as the first two SF triggers of Larson (1987), play a large part in grand design spirals, random collisions of interstellar clouds may provide the best explanation for dwarf galaxies with bursts of SF. Due to their small size, lack of strong spiral pattern, and sometimes solid-body rotation ({\\it e.g.,} Martimbeau \\etal 1994, Blok \\& McGaugh 1997), the star formation in dwarf galaxies is not triggered by compression due to gravitational density waves or by disk shear. Therefore, understanding SF in dwarf galaxies should be simpler than in other types of galaxies.  The characterization of the SF processes by a star formation rate (SFR) controlled by the interstellar gas density as a power law was first introduced by Schmidt (1959). The  volume density of young stars, $\\rho_*$, is related  to the volume density of HI gas in the Galactic disk as $\\rho_*=a \\; \\rho_{gas}^n$, where $n$ is a constant, probably $\\approx$2 for spiral galaxies. In other galaxies the convention is to express the quantities as projected densities of stars ($\\Sigma_*$) and of gas, as actually observed: $\\Sigma_*=A \\; \\Sigma_{gas}^n$. This is usually studied  by correlating the surface density of a young star tracer, such as the H$\\alpha$ surface brightness, with the gas column density.  The H$\\alpha$ emission from a galaxy measures its ongoing SFR (Kennicutt 1983). Gallagher \\etal (1984) derived an analytic relation between the detected H$\\alpha$ flux and the present SFR of a galaxy; similar relations were derived by Kennicutt \\etal (1994). The blue  luminosity of a galaxy, on the other hand,  measures its past star formation integrated over the last $\\sim 10^9$ yrs (Gallagher \\etal 1984). The newly formed stars, of which the more massive produce the Lyman continuum photons which ionize hydrogen and produce the H$\\alpha$ emission, contribute also to the blue light output of a galaxy. This contribution is minor in comparison to that from the stars already existing in a galaxy, unless the SF event is the first in the history of the galaxy or the star burst is unusually strong. Interestingly, Tresse \\& Maddox (1998) found recently that the H$\\alpha$ luminosity of a galaxy correlates with its blue absolute magnitude.  Kennicutt (1998) found that his parametrization of the Schmidt law fitted well the SF pattern of spiral and IR-selected starburst galaxies. An alternative to the Schmidt law, proposed by Silk (1997), fitted equally well. In this variant, the SFR per unit area scales with the ratio of the gas surface density to the local dynamical timescale: $\\Sigma_{SFR}\\propto\\frac{\\Sigma_{gas}}{\\tau_{dyn}} \\propto\\Sigma_{gas}\\times\\Omega_{gas}$, where $\\Omega_{gas}$ is the angular rotation speed and the scenario fits disk configurations.  Kennicutt (1998) adopted $\\Omega_{gas}=\\frac{V(R)}{R}$, where V(R) is the rotation velocity of the gas at a distance R.  Hunter \\etal (1998) tested a set of SF predictors on two small samples of dwarf galaxies, one measured by them and another derived from de Blok (1997). They found that the ratio of $\\Sigma_{gas}$  to the critical density for the appearance of ring instabilities did not correlate with the star formation, but that the stellar surface brightness did. From this, they concluded that possibly the stellar energy input provides the feedback mechanism for star formation.  We concentrate on a sample of late-type dwarf galaxies in the Virgo cluster (VC). The reason for selecting dwarfs was to limit the number of possible trigger mechanisms of SF; these objects are devoid of large-scale SF triggers, as explained above. Having only VC members limits the sample to a well-defined galaxy background; in addition, all objects are at $\\sim$the same distance and have HI information from the same source. We tested for correlations between the H$\\alpha$ emission and other observed quantities, in order to investigate mechanisms which regulate  SF in dwarf galaxies. The justification to correlate the H$\\alpha$ SFR index against $\\Sigma_{gas}$ is the finding of Kennicutt (1998) that a Schmidt-type law seems to fit large galaxies. If the SFR depends on the gas density ratioed to the dynamical timescale (Silk 1997, Kennicutt 1998), a correlation with $\\Sigma_{gas}\\times\\Omega_{gas}$ is expected. Finally, if the SFR depends on the local population of blue stars, as found by Hunter \\etal (1998), then a dependence on the average blue surface brightness is expected. We also tested the SFR against the ISM gas velocity, represented by the width of the HI line profile at 20\\% of its peak intensity ($\\sigma$(HI)), and against a combination of it and the surface density of HI, in a manner similar to that suggested by Silk (1997). ",
        "conclusions": "We tested correlations among parameters related to star formation, gas and stellar content, and internal dynamics on a sample of dwarf galaxies in the Virgo cluster. We found that both the Schmidt law and the more recent relation derived by Silk (1997) do not fit these galaxies as well as they do spirals (Kennicutt 1998). The strongest correlation  of the H$\\alpha$ surface brightness, which measures the present star formation strength, was with the  average blue surface brightness, supporting the proposition of Hunter \\etal (1998) that a feedback mechanism must be at work to regulate the present SF by the older stellar population."
    },
    "9804/astro-ph9804217_arXiv.txt": {
        "abstract": "Recent {\\it Hubble Space Telescope} observations have found that the horizontal branches (HB's) in the metal-rich globular clusters NGC~6388 and NGC~6441 slope upward with decreasing $\\bv$. Such a slope is not predicted by canonical HB models and cannot be produced by either a greater cluster age or enhanced mass loss along the red-giant branch (RGB). The peculiar HB morphology in these clusters may provide an important clue for understanding the second-parameter effect.  We have carried out extensive evolutionary calculations and numerical simulations in order to explore three non-canonical scenarios for explaining the sloped HB's in NGC~6388 and NGC~6441: i)~A high cluster helium abundance scenario, where the HB evolution is characterized by long blue loops; ii)~A rotation scenario, where internal rotation during the RGB phase increases the HB core mass; iii)~A helium-mixing scenario, where deep mixing on the RGB enhances the envelope helium abundance. All three of these scenarios predict sloped HB's with anomalously bright RR Lyrae variables. We compare this prediction with the properties of the two known RR Lyrae variables in NGC~6388. Additional observational tests of these scenarios are suggested. ",
        "introduction": "Many years ago Sandage \\& Wildey (1967) and van den Bergh (1967) discovered that the horizontal-branch (HB) morphology in the Galactic globular clusters (GC's) does not correlate tightly with the cluster metallicity. While theory and observation agree that the HB morphology on average becomes redder with increasing metallicity, there are many examples of GC's having very similar metallicities but markedly different HB morphologies, e.g., M3 versus M13, NGC~362 versus NGC~288. Thus some parameter besides metallicity must be affecting the evolution of the HB stars in these clusters. Indeed, theoretical HB models show that there are many possible ``second parameter\" candidates, e.g., the GC age, mass loss along the red-giant branch (RGB), helium abundance $Y$, alpha-element enhancement [$\\alpha$/Fe], and rotation (Fusi Pecci \\& Bellazzini 1998).  Recent {\\it Hubble Space Telescope} (HST) observations have revealed that the metal-rich GC's NGC~6388 and NGC~6441 contain a significant population of blue HB (BHB) stars and therefore exhibit a pronounced second-parameter effect (Piotto et al. 1997; Rich et al. 1997). From their analysis, Rich et al. concluded that neither a greater age nor enhanced RGB mass loss due to dynamical effects could satisfactorily account for the BHB populations in these GC's. Understanding the origin of these BHB stars has important implications for determining the formation history of the Galaxy and for interpreting the integrated spectra of the old metal-rich stellar populations in elliptical galaxies. If, for example, age were the second parameter in these GC's, then NGC~6388 and NGC~6441 would be among the oldest, if not the oldest, GC's in the Galaxy.  Sweigart \\& Catelan (1998) pointed out that the HB's of NGC~6388 and NGC~6441 have a pronounced upward slope with decreasing $\\bv$, with the mean luminosity at the top of the blue tail being nearly 0.5~mag brighter in $V$ than the well-populated red HB (RHB) clump. A similar upward slope is also present in the RHB itself (Piotto et al. 1997). As we shall see, such upward sloping HB's cannot be explained within the framework of canonical theory. More specifically, they cannot be attributed to differences in age or RGB mass loss---the two most prominent second-parameter candidates. We emphasize therefore that these GC's may provide a crucial clue for understanding the second-parameter effect.  In this {\\it Letter} we will investigate the following non-canonical scenarios for explaining the BHB populations and, in particular, the sloped HB's in NGC~6388 and NGC~6441: \\begin{enumerate} \\item {\\it High $Y$ scenario}, where a high helium abundance at the time of GC formation leads to unusually long blue loops during the HB evolution;  \\item {\\it Rotation scenario}, where internal rotation increases the HB core mass, making the HB both bluer and brighter;  \\item {\\it Helium-mixing scenario}, where deep mixing along the RGB increases the HB envelope helium abundance, also leading to bluer and brighter HB stars. \\end{enumerate}  \\noindent We will explore each of these possibilities under the simplest possible assumptions and will show that each can produce sloped HB's resembling those in NGC~6388 and NGC~6441.  We begin in Sec.~2 by discussing the failure of the canonical models to explain the observed HB's in NGC~6388 and NGC~6441; in Sec.~3, we describe each of the above non-canonical scenarios in more detail; in Sec.~4, we show that the RR Lyraes in NGC~6388 are brighter than field variables of similar [Fe/H], in agreement with these scenarios. Finally, in Sec.~5 we suggest some observational tests to discriminate among these scenarios. ",
        "conclusions": "The results in this {\\it Letter} demonstrate that canonical models fail to explain the sloped HB's in NGC~6388 and NGC~6441. This has the important implication that, at least for these GC's, the two most prominent second-parameter candidates---age and RGB mass loss---cannot account for the observed HB morphology. Simulations employing non-canonical models show that an increased initial GC helium abundance or a dispersion in the stellar rotational velocity (leading to a range in $M_{\\rm c}$ or $Y_{\\rm HB}$) may account for the observed HB morphology in these two GC's. Finally, we find that the RR Lyrae variables in NGC~6388 and 47~Tuc are substantially brighter than field variables of comparable [Fe/H]---in agreement with the non-canonical scenarios.  We can suggest a variety of observational tests for discriminating among the different non-canonical scenarios, including: \\begin{enumerate} \\item {\\it Analysis of R-ratios}: the $R$ ratios differ widely from one non-canonical scenario to the next, being largest for the high-$Y$ case (Figs.~2a,b). A low $R$ ratio (implying $Y_{\\rm MS} \\lesssim 0.35$) for NGC~6388 and NGC~6441 would argue against the high-$Y$ scenario;  \\item {\\it Deep photometry of NGC~6388 and NGC~6441}: this would determine whether the magnitude difference between the RR Lyraes and the turnoff, ``$\\Delta V$,\" is larger than expected for a ``47~Tuc-like\" age, as predicted by the non-canonical scenarios;  \\item {\\it RR Lyrae survey}: NGC~6388 and NGC~6441 should be surveyed for RR Lyraes, in order to discover new variables, especially closer to the cluster centers, and to define better the pulsation properties of the known RR Lyraes. In addition, we suggest a study of the RR Lyraes in the metal-rich GC's NGC~6304 (Hartwick, Barlow, \\& Hesser 1981; Layden 1995) and NGC~6652 (Ortolani, Bica, \\& Barbuy 1994);  \\item {\\it Surface gravities of blue HB stars}: $\\log\\,g$ measurements could discriminate among the different origins of the BHB stars in NGC~6388 and NGC~6441 (Crocker, Rood, \\& O'Connell 1988) and provide a strong constraint on whether the HB's are indeed brighter than predicted by the canonical models;  \\item {\\it High-resolution spectroscopy of red giants}: RGB stars in NGC~6388 and NGC~6441 should be searched for evidence of deep mixing (cf. Kraft 1994);  \\item {\\it Planetary Nebula (PN) in NGC~6441}: the PN recently discovered by Jacoby et al. (1997) in NGC~6441 may be an invaluable diagnostic tool for constraining the models. Helium is moderately enhanced over the solar value in this PN, while oxygen is more depleted than in the super-oxygen-poor stars which are found in some more metal-poor GC's and which are indicative of deep mixing (e.g., Kraft 1994). We speculate that this PN may be the progeny of a star that experienced helium mixing.  \\end{enumerate}"
    },
    "9804/astro-ph9804092_arXiv.txt": {
        "abstract": "We report the discovery of WN~J0717+4611, a highly polarized red quasar at $z=1.462$, selected from a sample of faint ($S_{1.4GHz} >$ 10 mJy) ultra steep spectrum ($\\alpha_{1400}^{327} < -1.3$; $S \\propto \\nu^{\\alpha}$) radio sources. VLA observations at 4.85 GHz show that the radio source is dominated by a 21 mJy core. Using the Keck telescopes, we obtained spectra with a combined wavelength range covering 4000~\\ang\\ to 1.6 \\micron, and spectropolarimetric observations between 4000~\\ang\\ and 9000~\\ang. We identify WN~J0717+4611 as a quasar based on broad emission of \\MgII~$\\lambda$2799 (FWHM $\\simeq$ 11,000 km s$^{-1}$), a stellar $R_S$-band Lick image and the high luminosity of the object ($M_B=-24.2$). The optical linear polarization is wavelength-independent at 15\\% in both the continuum and the emission lines, with the polarization angle orthogonal to the radio jet axis. We argue that the polarization mechanism is scattering of quasar light by surrounding small dust grains or electrons. WN~J0717+4611 appears highly reddened for a core-dominated quasar. ",
        "introduction": "The tight correlation in the Hubble $K - z$ diagram of the parent galaxies of luminous radio sources suggests that high redshift radio galaxies (HzRGs) may be used to study the formation and evolution of massive ellipticals (\\cite{lil89}, \\cite{eal97}, \\cite{wvb98}). HzRGs have now been found out to redshifts of $\\sim 4.5$, which is presumably close to the time they underwent their first burst of star formation. As a result of various vigorous search campaigns, $\\sim 20$ HzRG are now known at $z > 3$. Virtually all of these have been found by identifying sources with ultra-steep ($\\alpha \\lesssim -1$) radio spectra (\\cite{cha96}; \\cite{raw96}; \\cite{rot97}; \\cite{deb98a}). Although this ultra-steep spectrum (USS) search technique has now been proven the most effective way of finding HzRGs, no uniform large-area search could be done because the radio surveys used to construct the samples were either not sensitive enough, or had limited sky coverage. During the past years, several new, deep all-sky radio surveys have become available: the Faint Images of the Radio Sky at Twenty Centimeters at 1.4 GHz (FIRST, \\cite{bec95}), the Westerbork Northern Sky Survey at 325 MHz (WENSS, \\cite{ren97}), and the NRAO VLA Sky Survey at 1.4 GHz (NVSS, \\cite{con98}). These allow for the first time a major survey of faint radio sources with extreme spectral index properties.  We have used these new surveys to define a sample of USS sources, for further optical identification and spectroscopic studies (see \\cite{deb98a} for a summary). Our sample includes sources with radio flux densities $\\sim10 \\times$ fainter than previous large USS searches (\\cite{cha96}; \\cite{rot97}). Subsequent optical imaging and spectroscopy of our sample indicates that 2/3 of the USS sources with $R < 24$ are radio galaxies with redshifts $z > 2$. We will report in detail on these results in forthcoming papers.  In this paper, we present the discovery and follow-up observations of one of our USS sources, WN~J0717+4611 at $z = 1.462$. This source was selected for a special study because of its strong core-dominated FR II morphology (\\cite{far74}) and a bright stellar $R_S = 21.6$ identification. Such a combination of properties is unique in our USS sample, and its bright compact optical identification is reminiscent of 6C 1908+722, a $R = 21$ radio galaxy at $z = 3.515$ with a broad absorption line (BAL) system (\\cite{dey97}).  Our Keck optical spectropolarimetry observations do not reveal such a BAL system, but, in combination with the IR-spectroscopy that was originally obtained to study the age of the stellar population by means of the strength of the 4000~\\ang-break, it shows evidence for a highly polarized red quasar. We argue that the polarization is caused by scattering and not by synchrotron radiation, as in most highly polarized quasars (\\cite{coh97}). The spectral properties of the polarized light indicates that the red color of WN~J0717+4611 is strong evidence for dust obscuration, which has been one of the prime explanations for the red quasars (\\eg \\cite{web95}), although other mechanisms, like a host galaxy or synchrotron contribution, have also been proposed (\\eg \\cite{ben98}).  We describe the radio and optical imaging, IR spectroscopy, and optical spectropolarimetry in \\S\\ 2, and present the results in \\S\\ 3. We discuss the classification, the origin of the high polarization, and compare WN~J0717+4611 with possibly related objects in \\S\\ 4, and present our conclusions in \\S\\ 5. Throughout the paper, we use $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0$ and $\\Lambda = 0$. At $z = 1.462$, changing these values to $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0.1$, would reduce size and distance by 30\\%, which would not affect our conclusions. ",
        "conclusions": "Our main observational results for WN~J0717+4611 are:  1) The radio source has a core-dominated FR II morphology, with a total size of $\\sim 90$~kpc and an ultra-steep spectrum with an integrated spectral index $\\alpha_{1400}^{327} = -1.31$. The radio spectrum flattens with increasing frequency, presumably because of an increasing contribution from a flatter spectrum core.  2) The optical object is a quasar because it is stellar, has broad \\MgII~$\\lambda$2799 emission with (FWHM $\\simeq$ 11,000 km s$^{-1}$), and is very luminous with $M_B = -24.2$.  3) The object appears exceptionally reddened for its core-dominance of $R \\simeq 0.6$, compared to the statistical results of \\cite{bak95}. When dereddened with an $A_V = 1.2$, it matches well with the composite CSS quasar spectrum, which is already more reddened than the core-dominated quasar composite.  4) We detect a wavelength-independent 15\\% optical polarization, which is similar in both continuum and emission lines. The polarization angle is orthogonal to the radio axis, which, in the context of the unified models of quasars and radio galaxies, is consistent with scattering due to small dust grains or electrons located in the cone containing the radio jet.  We have compared WN~J0717+4611 with other HPQs and red quasars in the literature, and found no object that shares all observational characteristics. 5C1.195 (\\cite{wil98a}) seems to be a very similar object, and it was discovered in a similar way, indicating more of these objects may be found from radio-selected samples. The presence of a significant population of (highly) reddened quasars at cosmological redshifts is currently subject to considerable debate (\\cite{web95}; \\cite{ben98}). While our observations of a single object do not bear any statistical significance on the size of this population, they do show that highly reddened, dusty quasars do indeed exist."
    },
    "9804/astro-ph9804321_arXiv.txt": {
        "abstract": "We apply the pixel method of analysis (sometimes called ``pixel lensing'') to a small subset of the EROS-1\\ microlensing observations of the bar of the Large Magellanic Cloud (\\object{LMC}).  The pixel method is designed to find microlensing events of unresolved source stars and had heretofore been applied only to M31 where essentially all sources are unresolved. With our analysis optimised for the detection of long-duration microlensing events due to 0.01-1 $M_\\odot$ Machos, we detect no microlensing events and compute the corresponding detection efficiencies.  We show that the pixel method, applied to crowded fields, should detect 10 to 20 times more microlensing events for $M>0.05 M_\\odot$ Machos compared to a classical analysis of the same data which latter monitors only resolved stars.  In particular, we show that for a full halo of Machos in the mass range 0.1 -- 0.5 $M_\\odot$, a pixel analysis of the three-year EROS-1 data set covering $0.39\\,\\rm deg^2$ would yield $\\simeq 4$ events. ",
        "introduction": "Microlensing searches can probe the distribution of MAssive Compact Halo Objects (Machos) in the dark halo of our Galaxy or more distant galaxies (Paczy\\'nski (1986,1996), Griest (1991))\\nocite{Paczynski:1986,Griest:1991,Paczynski:1996}. Several lines of sight are now under investigation, and events have been claimed in several directions: towards the \\object{LMC} (Alcock et al.\\ 1997a, Renault et al.\\ 1997)\\nocite{Alcock:1997a,Renault:1997}, the SMC (Alcock et al.\\ 1997b, Pa\\-lan\\-que-Delabrouille et~al. 1998\\nocite{Alcock:1997b,Palanque-Delabrouille:1998}), in the direction of the Galactic Bulge (Alard et al.\\ 1997, Alcock et al. 1997c, Udalski et al 1994)\\nocite{Alcock:1997d,Alard:1997,Udalski:1994b} and more recently towards spiral arms (Derue et al.\\ 1998)\\nocite{Derue:1998}, giving some first evidences of the Macho distribution towards these lines of sight. These results are based on a star monitoring analysis: the fluxes of several millions of {\\it resolved} stars are monitored. As first discussed by Crotts (1992)\\nocite{Crotts:1992} and Baillon et al.\\ (1993), events due to {\\em unresolved} stars essentially escape detection of these analyses. Such stars, beyond the crowding limit or too dim to resolve, could significantly contribute to the number of detectable events.  This is illustrated by the detection of two variable objects in the MACHO analysis (Alcock et al. 1997a)\\nocite{Alcock:1997a} of \\object{LMC} data, which could not be resolved at their minimum luminosity. The detection of the variation was nevertheless possible because the reference images used to establish the catalogue of monitored stars were taken during their maximum luminosity, when the stars were resolved. However, such events can neither be kept for further considerations nor be included in the computation of detection efficiencies, in the star monitoring approach.  In this paper, we apply a pixel analysis to the EROS~91-92 data (10\\% of the whole EROS-1 CCD data set) of the \\object{LMC Bar}. The idea is to monitor the flux of all the pixels present on the images, thus achieving a good sensitivity to the whole stellar content of the images. The magnification of one unresolved star can be detected as a variation of the pixel flux, provided that the magnification is high enough. -- In the following, we will refer to this approach as pixel monitoring, as opposed to star monitoring referring to classical analyses restricted to resolved stars. -- The main uncertainties of this approach concern our ability to account properly for variations of the observational conditions, and to be able to disentangle intrinsic variations from observational systematics.  In Paper~I (Melchior et al. 1998a)\\nocite{Melchior:1998a}, dedicated to the description of the treatment of the data and the production of $2.1 \\times 10^6$ super-pixel light curves, we have shown that an average stability of $1.8\\%$ of the super-pixel flux is achieved in blue and $1.3\\%$ in red, about twice the expected photon noise.  This homogeneous set of super-pixel light curves is called AGAPEROS: each of these light curves covers a period of $120$ days and is composed of about 90 measurements. With this rather short period of observation, we show here how it is possible to investigate the Macho mass range of interest ($M\\simeq 10^{-2} - 1 M_{\\odot}$) with the existing EROS-1 CCD data set, initially designed for short time scale events in a mass range ($M\\simeq 10^{-8} - 10^{-3} M_{\\odot}$) where no event has been detected.  \\subsection*{Microlensing selection with the Pixel Method} \\label{section:pixmeth} We first present the simple formalism used to describe the pixel events in which we are interested.  The pixel flux $\\phi^{p}$, affected by a microlensing event, can be written as: \\begin{equation} \\phi^{p}_i = \\alpha_i A(t_i) \\phi^*_i + \\phi^{\\rm bg}_i \\label{eq:phi} \\end{equation} where $i$ is the measurement number, $t$ is the time, $\\alpha$ is the seeing fraction, $A$ the magnification, $\\phi^*$ the flux of the star of interest at rest (i.e., unmagnified), and $\\phi^{\\rm bg}$ includes the sky and stellar backgrounds. The magnification (Paczy\\'nski 1986)\\nocite{Paczynski:1986} depends on the normalized impact parameter $u(t_i)$: \\begin{equation} A(t_i) = \\frac{u^2+2}{u\\sqrt{u^2+4}} \\end{equation} with \\begin{eqnarray} u(t_i) = \\sqrt{u_0^2 + {\\left(\\frac{v_T}{R_E}\\right)}^2 (t_i - t_0)^2} \\nonumber \\end{eqnarray} where $v_T$ is the Macho transverse velocity, $u_0$ and $t_0$ the impact parameter and time of closest approach, and $R_E$ the Einstein radius: \\begin{eqnarray} R_E = \\sqrt{\\frac{4 G M}{c^2} \\frac{D_{OL} D_{LS}}{D_{OS}}}\\nonumber \\end{eqnarray} The typical time scale of the variation is the Einstein radius crossing time: \\begin{equation} t_E = \\frac{R_E}{v_T} \\end{equation}  In Paper~I, we showed that the variations of the observational conditions, which obviously affect Eq.~\\ref{eq:phi}, can be corrected: each image is first geometrically aligned with the reference image. Then sky background and absorption factor are corrected to the values of the reference image.  Finally, each super-pixel flux is adjusted to take account of the seeing variation, affecting $\\alpha_i$.  Since the mean seeing is about $3$ arc-second, we estimate $\\alpha$ to be on average $0.7$ for the corrected $3.\\hskip-2pt ''6 \\times 3.\\hskip-2pt ''6$ super-pixel light curves, obtained in Paper~I.  We showed that in the absence of microlensing events ($A\\left( t_i \\right) = 1$), we achieve a proper understanding of the errors affecting these light curves, and that Eq.~\\ref{eq:phi} describes to a good approximation the light curves we are studying.  The usual requirements used for the selection of microlensing events detected by star monitoring can be applied here: \\begin{itemize} \\item As the microlensing phenomenon is a transient and rare phenomenon, it should produce a {\\em unique} significant variation in the star flux. \\item It must be {\\em achromatic}. This characteristic has two applications for a pixel analysis: the time of the maximum has to be the same in both colours and the ratio \\begin{equation} \\frac{ {{\\phi^p_i}^{(B)}} - \\phi^{(B)}_{\\rm bg}}{{\\phi^p_i}^{(R)} - \\phi^{(R)}_{\\rm bg}} \\label{eq:ach} \\end{equation} must remain constant, during the variation. \\item The {\\em shape} must be compatible with Eq.~\\ref{eq:phi}. \\end{itemize} The first criterion allows us to remove recurrent variable stars as well as most of the noise, while the two other criteria will be applied to the few remaining light curves at the end of the selection process. \\begin{itemize} \\item Last, we consider the fact that the probability for a star to be lensed is independent of its type. This will allow us to reject specific populations of variable stars. \\end{itemize}  In Sect. \\ref{section:simul}, we present the Monte-Carlo simulations used in this article. Then, in Sect. \\ref{section:trigger}, we describe step by step the selection procedure designed to detect microlensing events and applied to these light curves. In Sect. \\ref{section:colmag}, we show how the selected variations can be eliminated as compatible with variable stars. In Sect. \\ref{section:detect}, we discuss the detection efficiencies achieved by this analysis, and compare the number of expected events with the sensitivity of star monitoring. We rely on these results in Sect. \\ref{section:perspec} to discuss the possible prospects of this approach. ",
        "conclusions": "While the EROS-1 CCD data set has already excluded the small-mass Macho range, we have shown that it is possible with a pixel analysis of the same data set to probe the mass range ($0.01 - 1. M_\\odot$) where all the known events have been detected.  We thus demonstrate for the first time the efficacy of the pixel method for \\object{LMC} images.  We have shown with the computation of detection efficiencies that the gain in detectable microlensing events with pixel monitoring is significant: {\\em pixel monitoring can detect $15$ times more microlensing events than star monitoring} for very crowded LMC bar fields.  With simple selection criteria we have been able to show that no microlensing events are present in this data set. The criteria imposed mainly rely on the uniqueness of the variation, on the achromaticity and on the fact that the microlensing phenomenon -- a geometrical effect -- should affect all kind of stars independently of their type. This demonstrates that the noise which affects super-pixel light curves can be reduced to a level that is adequate to conduct the analysis.  Our analysis, optimised for the detection of long duration events, has only detected variable stars, and we have shown our ability to reject such variable stars with the CMD. We have then demonstrated the efficacy of this approach for the detection of variable stars, which remain the main background for microlensing.  A comprehensive analysis of the variable stars rejected by this analysis is in progress (Melchior et al. 1998b)\\nocite{Melchior:1998c}.  Such an analysis can detect microlensing magnification of unresolved stars up to $V=24$ and is thus complementary to star monitoring. On-line pixel analysis and follow-up will probably be necessary to better discriminate possible variable stars and to achieve a higher quality photometry.  The short period of observation analysed here -- $120$ days to be compared to $3$ years for EROS-1 plates and $2$ years for MACHO -- and the relatively small field -- $100$ times smaller than the one analysed by the EROS group (plates) (Ansari et al. 1996)\\nocite{Ansari:1996} and $60$ times smaller than the field analysed by the MACHO group (Alcock et al.\\ 1997a)\\nocite{Alcock:1997a} -- explain the relatively small number of events expected. For the first time, we are able to provide detection efficiencies for a pixel analysis. These results allow us to estimate that a pixel analysis of the complete existing EROS-1 CCD data base could detect about $\\simeq 4. \\times f$ microlensing events in the mass range of interest ($0.1 - 0.5 M_{\\odot}$)... with a field of only $0.39\\,\\rm deg^2$!"
    },
    "9804/astro-ph9804117_arXiv.txt": {
        "abstract": "The first sources of ionizing radiation to condense out of the dark and neutral IGM sent ionization fronts sweeping outward through their surroundings, overtaking other primordial gas-clouds and photoevaporating them. Results are presented of the first gas dynamical simulations of this process, including radiative transfer, along with some observational diagnostics. ",
        "introduction": "The r\\^ole of hydrogen molecules as cooling agents in primordial gas clouds, necessary to form the first stars out of a gas of H and He with no heavier elements, is intimately related to the fate of those clouds in the presence of ionizing and dissociating radiation. That such radiation had an important effect on these primordial clouds is certain. The neutral, opaque IGM out of which the first bound objects condensed was dramatically reheated and reionized at some time between a redshift $z\\approx50$ and $z\\approx5$ by the radiation released by some of these objects (cf. [1] and references therein). When the first sources turned on, they ionized their surroundings by propagating weak, R-type ionization fronts which moved outward supersonically with respect to both the neutral gas ahead of and the ionized gas behind the front, racing ahead of the hydrodynamical response of the IGM [2,3]. The effect of density inhomogeneity on the rate of I-front propagation was previously described by a mean ``clumping factor'' $c_l>1$, which slowed the I-fronts by increasing the average recombination rate per H atom inside clumps. This suffices to describe the rate of I-front propagation as long as the clumps are either not self-shielding or, if so, only absorb a fraction of the ionizing photons emitted by the central source. What is the dynamical effect of the I-front on the density inhomogeneity it encounters, however?  The answer depends on the size and density of the clumps overtaken by the I-front. The fate of linear density fluctuations depends upon their Jeans number, $L_J\\equiv\\lambda/\\lambda_J$, the wavelength in units of the baryon Jeans length in the IGM at temperatures of order $10^4\\rm K$. Fluctuations with $L_J<1$ find their growth halted and reversed (cf. [4]). For nonlinear density fluctuations, however, the answer is more complicated, depending upon at least three dimensionless parameters, their internal Jeans number, $L_J\\equiv R_c/\\lambda_J$, the ratio of the cloud radius $R_c$ to the Jeans length $\\lambda_J$ inside the cloud at about $10^4{\\rm K}$, their ``Str\\\"omgren number'' $L_s\\equiv R_c/\\ell_s$, the ratio of the cloud radius $R_c$ to the Str\\\"omgren length $\\ell_s$ inside the cloud (the length of a column of gas within which the unshielded arrival rate of ionizing photons just balances the total recombination rate), and their optical depth to H ionizing photons at 13.6 eV, $\\tau_{\\rm H}$, before ionization. If $\\tau_{\\rm H}<1$, the I-front sweeps across the cloud, leaving an ionized gas at higher pressure than its surroundings, and exits before any mass motion occurs in response, causing the cloud to blow apart. If $\\tau_{\\rm H}>1$ and $L_s>1$, however, the cloud shields itself against ionizing photons, trapping the I-front which enters the cloud, causing it to decelerate inside the cloud to the sound speed of the ionized gas before it can exit the other side, thereby transforming itself into a weak, D-type front preceded by a shock. Typically, the side facing the source expels a supersonic wind backwards towards the source, which shocks the IGM outside the cloud, while the remaining neutral cloud material is accelerated away from the source by the so-called ``rocket effect'' as the cloud photoevaporates (cf. [5]). As long as $L_J<1$ (the case for gas bound to dark halos with virial velocity less than $\\rm 10\\,km\\,s^{-1}$), this photoevaporation proceeds unimpeded by gravity. For halos with higher virial velocity, however, $L_J>1$, and gravity competes with pressure forces. For a uniform gas of H density $n_{\\rm H,c}$, located a distance $r_{\\rm Mpc}$ (in Mpc) from a UV source emitting $N_{\\rm ph,56}$ ionizing photons (in units of $\\rm10^{56}s^{-1}$), the Str\\\"omgren length is only $\\ell_s\\cong(50\\,{\\rm pc}) (N_{\\rm ph,56}/r_{\\rm Mpc}^2)(n_{\\rm H,c}/0.1\\,\\rm cm^{-3})^{-2}$. We focus in what follows on the self-shielded case which traps the I-front. Some of these results were summarized previously by us in [6]. ",
        "conclusions": ""
    },
    "9804/astro-ph9804051_arXiv.txt": {
        "abstract": "We extend the formalism of the relativistic thermal Sunyaev-Zel'dovich effect to the system moving with a velocity $\\vec{\\beta} \\equiv \\vec{v}/c$ with respect to the cosmic microwave background radiation.  In the present formalism, the kinematic Sunyaev-Zel'dovich effect for the cluster of galaxies with a peculiar velocity $\\vec{\\beta}$ is derived in a straightforward manner by the Lorentz boost of the generalized Kompaneets equation.  We give an analytic expression for the kinematic Sunyaev-Zel'dovich effect which is valid up to O($\\beta^{2}$) with the power series expansion approximation in terms of $\\theta_{e} \\equiv k_{B} T_{e}/mc^{2}$, where $T_{e}$ and $m$ are the electron temperature and the electron mass, respectively.  It is found that the relativistic corrections to the kinematic Sunyaev-Zel'dovich effect are significant.  For a typical electron temperature $k_{B} T_{e}$ = 10keV, one obtains $-8.2\\%$ and $+1.3\\%$ corrections from the $O(\\beta \\theta_{e})$ and $O(\\beta \\theta_{e}^{2})$ contributions, respectively.  The $O(\\beta^{2})$ correction is extremely small, $+0.2\\%$ for $\\beta=1/300$ at $k_{B} T_{e}$ = 10keV.  Therefore it can be safely neglected.  These relativistic corrections are directly reflected on the determination of the peculiar velocity $\\beta$ of the cluster of galaxies with the observation of the kinematic Sunyaev-Zel'dovich effect. ",
        "introduction": "Compton scattering of the cosmic microwave background radiation (CMBR) by hot intracluster gas --- the thermal Sunyaev-Zel'dovich effect --- (Zel'dovich \\& Sunyaev 1969; Sunyaev \\& Zel'dovich 1972; Sunyaev \\& Zel'dovich 1980a) provides a useful method to measure the Hubble constant $H_{0}$ (Gunn 1978; Silk \\& White 1978; Birkinshaw 1979; Cavaliere, Danese, \\& De Zotti 1979; Birkinshaw, Hughes, \\& Arnaud 1991; Birkinshaw \\& Hughes 1994; Myers et al. 1995; Herbig et al. 1995; Jones 1995; Markevitch et al. 1996; Holzapfel et al. 1997a; Holzapfel et al. 1997b).  The Sunyaev-Zel'dovich formula has been derived from a kinetic equation for the photon distribution function taking into account the Compton scattering by electrons --- the Kompaneets equation --- (Kompaneets 1957; Weymann 1965).  The Kompaneets equation has been derived with the non-relativistic approximation for the electron.  However, the electrons in the clusters of galaxies are extremely hot, $k_{B} T_{e}$ = $5 \\sim 15$keV (Arnaud et al. 1994; Markevitch et al. 1994; Markevitch et al. 1996; Holzapfel et al. 1997a; Holzapfel et al. 1997b).  Recently attempts have been made to include the relativistic corrections in the Sunyaev-Zel'dovich effect (Rephaeli 1995; Rephaeli \\& Yankovitch 1997).  However, it appears that the calculations have not been carried out in a manifestly covariant form.  For example, equation (4) in Rephaeli (1995), which comes from Chandrasekhar (1950), is a non-relativistic formula.  Since the extension of the Kompaneets equation to the relativistic regime is extremely important in view of many recent measurements of the Hubble constant $H_{0}$ with the use of the Sunyaev-Zel'dovich effect, we will solve the kinetic equation for the photon distribution function in a manifestly covariant form taking into account the Compton scattering by electrons.  Very recently a generalized Kompaneets equation has been derived by several groups: Stebbins (1997); Challinor \\& Lasenby (1998); Itoh, Kohyama \\& Nozawa (1998).  By using the generalized Kompaneets equation, analytic expressions for the Sunyaev-Zel'dovich effect have been derived as a power series of $\\theta_{e}$ = $k_{B}T_{e}/mc^{2}$, where $T_{e}$ and $m$ are the electron temperature and the electron mass, respectively.  It has been shown that the results obtained by the power series expansion agree with the previous numerical calculations by Rephaeli (1995) and Rephaeli \\& Yankovitch (1997), thereby proving the validity of their method.  It is really impressive to see their result agreeing with the result of the manifestly covariant calculation.  In particular, the convergence of the power series expansion has been carefully studied in the paper by Itoh, Kohyama \\& Nozawa (1998), where the analytic expressions up to $O(\\theta_{e}^{5})$ have been derived and the results have been compared with those of the direct numerical integration of the Boltzmann equation.  It has been shown that the power series expansion approximation is sufficiently accurate for the region $k_{B}T_{e} \\leq 15$keV by taking into account up to $O(\\theta_{e}^{5})$ contributions.  On the other hand, Sunyaev \\& Zel'dovich (1980b) have predicted that the CMBR spectrum is also distorted if the cluster of galaxies has a peculiar velocity with respect to the CMBR.  This effect is called the kinematic Sunyaev-Zel'dovich effect.  Rephaeli \\& Lahav (1991) and Haehnelt \\& Tegmark (1996) have discussed the possibility of the determination of the magnitude of the peculiar velocity by measuring the kinematic Sunyaev-Zel'dovich effect.  Observation of the kinematic Sunyaev-Zel'dovich effect has been done by Holzapfel et al. (1997c) and the limits on the peculiar velocities have been given for the clusters Abell 2163 and Abell 1689.  At the present stage the error bars are too large to give definite magnitudes of the peculiar velocities with the kinematic Sunyaev-Zel'dovich effect.  However, more accurate future observations are promising in the determination of the peculiar velocities.  Therefore, relativistic corrections should be carefully taken into account in order to make the theoretical expressions for the kinematic Sunyaev-Zel'dovich effect as accurate as that for the thermal Sunyaev-Zel'dovich effect.  The purpose of the present paper is to calculate the relativistic correction to the kinematic Sunyaev-Zel'dovich effect.  In the present case, there exist two kinds of the relativistic corrections: one is due to the relativisitc electron temperature parameterized by $\\theta_{e}$, and the other is due to the relativistic peculiar velocity parameterized by $\\beta \\equiv v/c$.  In particular we are interested in the interference contributions such as $O(\\beta \\theta_{e})$, $O(\\beta \\theta_{e}^{2})$ and $O(\\beta^{2})$.  In the present paper we will derive the analytic expressions for the thermal and kinematic Sunyaev-Zel'dovich effects from a generalized Kompaneets equation by applying the Lorentz boost to the direction of the peculiar velocity of the cluster of galaxies.  The present paper is organized as follows.  In $\\S$ 2, the general formalism will be presented in deriving the thermal and kinematic Sunyaev-Zel'dovich effects on the same footing.  With the power series expansion approximation, an analytic expression including the relativistic effects will be derived for the thermal and kinematic Sunyaev-Zel'dovich effects.  Numerical results will be presented in $\\S$ 3.  Finally  concluding remarks will be given in $\\S$ 4. ",
        "conclusions": ""
    },
    "9804/astro-ph9804098_arXiv.txt": {
        "abstract": "We report HST NICMOS observations of the central region of NGC 5128 at 2.2\\mic\\ and in \\PA. The continuum images show extended emission typical of an elliptical galaxy and a strong unresolved central source we identify  as the nucleus of the galaxy. Its position is consistent with ground-based IR and radio data, and with the peak of reddening found with WF/PC-1. In \\PA, we detect a prominent elongated structure, centered on the nucleus, extended by $\\simeq2$\\arcsec at a position angle of $\\simeq$33\\arcdeg, and with a major to minor axis ratio of $\\sim2$. We interpret this as an inclined, $\\sim 40$ parsec diameter, thin nuclear disk of ionized gas rather than  a jet-gas cloud interaction. We do see several weaker \\PA\\ features, some of which may be circumnuclear gas clouds shocked by the X-ray/radio jet. The disk is one of the smallest ever observed at the nucleus of an AGN. It is not perpendicular to the jet, but consistent with being oriented along the major axis of the bulge.  If it represents the warped outer portion of an accretion disk around a black hole, we conclude that even on the scale of a few parsecs, the disk is dominated by the galaxy gravitational potential and not directly related to the symmetry axis of the AGN. ",
        "introduction": "NGC 5128 (Centaurus A), the nearby giant elliptical galaxy, hosts the closest powerful active galactic nucleus.  This double-lobed radio source contains a strong jet discovered in X-rays (\\cite{schreier:79}, \\cite{feigelson:81}) and well-studied in radio with the VLA (\\cite{schreier:81}, \\cite{burns:83}, \\cite{clarke:86}) and with the VLBI (cf. \\cite{jones:96}). The proximity of NGC 5128 makes it one of the best candidates for studying the inner region around a massive black hole, assumed by the standard model to be at the core of all AGNs (e.g. \\cite{blandford:91}, \\cite{antonucci:93}). However, the large warped dust lane, approximately an arcminute wide with associated gas, young stars, and HII regions, dominates the morphology and polarization properties at visible wavelengths, effectively obscuring the nucleus and inner half-kiloparsec of the galaxy and the jet.  Ground-based IR observations have provided evidence for a strong, heavily polarized source at the nucleus (cf. \\cite{bailey:86}, \\cite{packham:96}). Previous HST WF/PC-1 imaging polarimetry of the inner region of NGC5128 (\\cite{schreier:96}) identified the nucleus of the galaxy with the most obscured and reddened feature of the emission, several arcseconds SW of the brightest optical emission near the center of the dust lane and consistent with the IR results.  While the discovery of gas disks at the centers of nearby galaxies with HST has opened up new possibilities for determining masses of central black holes by studying kinematics of the disks (cf. \\cite{harms:94}, \\cite{ferrarese:96}, \\cite{macchetto:97},  \\cite{bower:98}) the dust lane does not permit such studies for NGC 5128 in the optical.  We report here HST Near Infrared Camera and Multi-Object Spectrometer (NICMOS) observations at 2.2\\mic\\ and \\PA\\ which allow us, for the first time, to study the structure of the circumnuclear region at high resolution.  We summarize the observations and data reduction in Section 2, the results in Section 3, and discuss the implications of these data for standard AGN and jet models in Sections 4 and 5.  Throughout, we assume a distance to Centaurus A of 3.5 Mpc (\\cite{hui:93}), whence 1\\arcsec$\\simeq$17pc. ",
        "conclusions": "We summarize the key results of our 2.2 $\\mu$m continuum and \\PA\\ observations of the inner region of NGC 5128 as follows:  1) We see an extended continuum source with the rather regular and smooth structure expected for an elliptical galaxy. This smoothness contrasts strongly with the filamentary structure observed in the WF/PC I-band image.  2) We identify a strong unresolved ($r<$0\\farcs1) central source as the nucleus of the galaxy.  Its position is consistent with ground-based IR and radio data and the peak of reddening and polarization found with WF/PC-1.  The nucleus has an intensity of (2.3$\\pm$0.1)\\xten{-15} erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$ based on our continuum observations, and a flux of (5.7$\\pm$0.6)\\xten{-14} erg cm$^{-2}$ s$^{-1}$ in \\PA.  3) We see a prominent elongated structure in \\PA\\, centered on the nucleus, with a major to minor axis ratio of $\\sim2$, extending $\\simeq2$\\arcsec along position angle $\\simeq$ 33\\arcdeg, perpendicular to the dust lane.  We interpret this as an inclined $\\sim 40$ parsec diameter thin nuclear disk of ionized gas.  4) We see several smaller \\PA\\ features, including three relatively strong ones, which appear spatially related to the X-ray/radio jet.  We interpret them as circumnuclear gas clouds shocked by the jet.  The disk, with a radius of $\\sim 1\\farcs2$, corresponding to $\\sim20$pc at 3.5 Mpc, is among the smallest ever observed at the nucleus of an AGN in the optical/near-IR.  It is significantly smaller than the hundred or more parsec stellar disks presumed to exist at the centers of many galaxies.  Even at this small scale, the disk is not perpendicular to the jet. It is, however, consistent with being perpendicular to the dust lane and oriented along the major axis of the bulge. If it represents the warped outer portion of an accretion disk around the active nucleus, we conclude that even a few parsecs from the nucleus, the disk is dominated by the galaxy gravitational potential and not by the symmetry of the AGN and its jet.  Further NICMOS observations are planned in 2$\\mu$m polarized light and [Fe\\,II] to provide emission mechanism diagnostics and determine the geometry of the nuclear radiation field. If a warped disk illuminated by the nucleus is the correct model, systematic changes in the polarization angle as a function of disk position should be seen.  High spatial resolution infrared spectroscopy is planned to measure gas kinematics in the disk, and thus determine the mass of the black hole."
    },
    "9804/astro-ph9804179_arXiv.txt": {
        "abstract": "Six high-resolution ($\\lambda/\\Delta\\lambda\\approx 3000$) far-UV ($\\lambda\\lambda=910$--1210~\\AA ) spectra of the magnetic cataclysmic variable AM~Herculis were acquired in 1996 November during the flight of the ORFEUS-SPAS II mission. AM~Her was in a high optical state at the time of the observations, and the spectra reveal emission lines of O~VI $\\lambda\\lambda 1032$, 1038, C~III $\\lambda 977$, $\\lambda 1176$, and He~II $\\lambda 1085$ superposed on a nearly flat continuum. Continuum flux variations can be described as per G\\\"ansicke et~al.\\ by a $\\approx 20$ kK white dwarf with a $\\approx 37$ kK hot spot covering a fraction $f\\sim 0.15$ of the surface of the white dwarf, but we caution that the expected Lyman absorption lines are not detected. The O~VI emission lines have broad and narrow component structure similar to that of the optical emission lines, the C~III and He~II emission lines are dominated by the broad component, and the radial velocities of these lines are consistent with an origin of the narrow and broad components in the irradiated face of the secondary and the accretion funnel, respectively. The density of the narrow- and broad-line regions is $n_{\\rm nlr}\\sim 3\\times 10^{10}~\\rm cm^{-3}$ and $n_{\\rm blr}\\sim 1\\times 10^{12}~\\rm cm^{-3}$, respectively, yet the narrow-line region is optically thick in the O~VI line and the broad-line region is optically thin; apparently, the velocity shear in the broad-line region allows the O~VI photons to escape, rendering the gas effectively optically thin. Unexplained are the orbital phase variations of the emission-line fluxes. ",
        "introduction": "AM~Herculis is the prototype of a class of strongly magnetic ($B\\approx 10$--100 MG) cataclysmic variables (see \\cite{cro90} for a review and the volume by \\cite{buc95} for recent results). The strong field locks the white dwarf into corotation and the accreting matter is channeled along the field lines for most of its trajectory from the secondary's inner Lagrange point to the white dwarf magnetic pole. A strong shock forms in the stream of freely falling matter close to the white dwarf surface, thermalizing a significant fraction of its kinetic energy and raising its temperature to a few tens of keV. The accretion region is therefore a source of hard X-rays, but soft X-rays are also produced by reprocessing of the hard X-rays in the white dwarf photosphere and by (with a tip `o the hat to Dr.\\ Bengue) ``deep heating'' of the white dwarf photosphere by blobs of material which penetrate to large optical depths before thermalizing their kinetic energy (\\cite{kui82}). The dominant source of the UV continuum is the the white dwarf photosphere, while the optical continuum is produced by free-free and bound-free emission from the funnel above the accretion shock and possibly the irradiated face of the secondary (\\cite{sch91}). Sources of IR through FUV emission lines include the irradiated face of the secondary, the accretion stream, and the accretion funnel. The contribution of each of these emission regions to the net spectrum can be distinguished by the phasing and radial velocity amplitudes of the emission lines, and the phasing and amplitude of continuum flux variations. Because of the small size and high X-ray luminosities of these binaries, photoionization of all of these emission regions is likely to be important, with ionization parameters $\\xi\\equiv L/nr^2 =100\\, L_{33}/n_{11}r_{10}^2$, where $L_{33}$ is the X-ray luminosity in units of $10^{33}~\\rm ergs~s^{-1}$, $n_{11}$ is the particle density in units of $10^{11}~\\rm cm^{-3}$, and $r_{10}$ is the distance from the X-ray emission region in units of $10^{10}$~cm.  AM~Her has been observed in the UV numerous times with {\\it IUE\\/} (\\cite{ray79}; \\cite{hei88}; \\cite{gan95}; \\cite{sil96}; \\cite{gan98}) but on only two previous occasions in the FUV: once in 1993 September during the flight of the ORFEUS-SPAS I mission (\\cite{ray95}) and once in 1995 March during the Astro-2 mission (\\cite{gre96}). We here discuss six FUV spectra of AM~Her acquired in 1996 November during the ORFEUS-SPAS II mission. These observations are superior to those of ORFEUS~I because of the extensive phase coverage and, thanks to the $\\approx 100$\\% improvement in the effective area (\\cite{hur96}; \\cite{hur97}), the higher signal-to-noise ratio of the individual spectra. Combined with the results from recent EUV and X-ray observations of AM~Her, the ORFEUS~II spectra allow a detailed picture to be drawn of the locations of, and physical conditions in, the various FUV emission regions of this compact magnetic binary. ",
        "conclusions": ""
    },
    "9804/gr-qc9804078_arXiv.txt": {
        "abstract": "s{ String cosmology models predict a stochastic cosmic background of gravitational waves with a characteristic spectrum. I describe the background, present astrophysical and cosmological bounds on it, and outline how it may be possible to detect it with gravitational wave detectors. } ",
        "introduction": "A robust prediction of models of string cosmology which realize the pre-big-bang scenario \\cite{sfd,pbb} is that our present-day universe contains a cosmic gravitational wave background \\cite{bggv,gg,gwg},  with a spectrum which is quite different than that predicted by many other early-universe  cosmological models \\cite{grishchuk,turner,myreview,maggiore}.  In the pre-big-bang scenario the evolution of the universe starts from a state of very small curvature and coupling and then undergoes a long phase of dilaton-driven kinetic inflation reaching roughly Planckian energy densities \\cite{exit1},  and at some later time joins smoothly standard  radiation dominated cosmological evolution, thus giving rise to a singularity free inflationary cosmology.  It is the purpose of this talk to explain the properties of the cosmic gravitational  wave background predicted by string cosmology, to explain astrophysical and cosmological bounds  on its shape and strength and to show that currently operating and planned gravitational  wave detectors could further constrain the spectrum and perhaps even detect it. The emphasis  is on the general properties of the spectrum and   the basic experimental setup needed for its detection, rather than on precise numerical estimates and technical details.  Because the gravitational interaction is so weak, a background of gravitational radiation decouples from the matter in the universe  at very early times and carries with it a picture of the state of the universe  when energy densities and temperatures were extreme. The weakness of the gravitational interaction makes a detection of such a background very hard, and necessitates a strong signal. String cosmology provides perhaps the strongest source possible: the whole universe, accelerated to roughly Planckian energy densities. A discovery of any primordial gravitational  wave background, and in particular, the one predicted by string cosmology,  could therefore provide unrivaled exciting information on the very early universe. ",
        "conclusions": ""
    },
    "9804/astro-ph9804209_arXiv.txt": {
        "abstract": "High resolution 21 cm observations of the Ursa Major cirrus revealed highly filamentary structures down to the 0.03 pc resolution. These filaments, still present in the line centroid map, show multi-Gaussian components and seem to be associated with high vorticity regions. Probability density functions of line centroid increments and structure functions were computed on the line centroid field, providing strong evidences for the presence of turbulence in the atomic gas. ",
        "introduction": "Many statistical studies of the density and velocity structure of dense interstellar matter have been done on molecular clouds where turbulence is seen as a significant support against gravitational collapse that leads to star formation. Less attention has been devoted to turbulence in the Galactic atomic gas (HI). The cold atomic component ($T \\sim 100$ K, $n \\sim 100$ cm$^{-3}$), alike molecular gas, is characterized by multiscale self-similar structures and non-thermal linewidths.  A detailed and quantitative study of the turbulence and kinematics of HI clouds has never been done. Here we present a preliminary analysis of this kind based on high resolution 21 cm observations of an HI cloud located in the Ursa Major constellation. To characterize the turbulent state of the atomic gas, a statistical analysis of the line centroid field has been done. We have computed probability density functions of line centroid increments and structure functions. ",
        "conclusions": "The previous statistical analysis supports the presence of turbulence in the Ursa Major cirrus. The various observational results shown here (the presence of clearly non-Gaussian wings in PDFs of line centroid velocity increments and the slopes of the structure functions) are recognized as signatures of turbulence in terrestrial experiments or numerical simulations. On the other hand, astrophysical observations lack three dimensions in phase space; any detailed comparison with numerical simulations or laboratory results must be done with great caution. Here the 3D geometry of the cloud is unknown, it is thus premature to interpret further the various SFs' exponents.  Numerical simulations and laboratory experiments of turbulent flows are also characterized by filamentary vortical structures and dissipative sheets around them \\linebreak (\\cite{sreenivasan97}). These highly localized dissipative regions may play an important role in the thermodynamics and chemical evolution of the interstellar medium (\\cite{falgarone95}). Identifying such regions would be of considerable astrophysical interest. The present study of high resolution 21 cm observations reveals filamentary HI structures characterized by doubled-peak spectra and minor axis velocity gradient. These features may be interpreted as a rotation around the major axis of the filament. Furthermore, the main filaments observed here are associated with high vorticity zones traced by the PDF of velocity increments.  The filaments we have observed in the Ursa cloud cannot be the coherent vortices which develop at scales close to the dissipation scale of turbulence because the linear resolution of our observations (0.03 pc) is more than 100 times larger than the dissipation scale in diffuse clouds (see Falgarone, these proceedings). Instead, we may be observing a larger structure, made of very small vortices, braided together, as the result of their interaction and merging. This merging of small vortices into larger vortices has been followed in the numerical simulations of Vincent and Meneguzzi (1994) and Porter et al. (1994). This would explain why the velocity pattern is more complex than a simple rotation."
    },
    "9804/astro-ph9804045_arXiv.txt": {
        "abstract": "Extended source size effects have been detected in photometric monitoring of gravitational microlensing events. We study similar effects in the centroid motion of an extended source lensed by a point mass. We show that the centroid motion of a source with uniform surface brightness can be obtained analytically. For a source with circularly symmetric limb-darkening profile, the centroid motion can be expressed as a one-dimensional integral, which can be evaluated numerically. We found that when the impact parameter is comparable to the source radius, the centroid motion is significantly modified by the finite source size. In particular, when the impact parameter is smaller than the source radius, the trajectories become clover-leaf like. Such astrometric motions can be detected using space interferometers such as the Space Interferometry Mission. Such measurements offer exciting possibilities of determining stellar parameters such as stellar radius to excellent accuracy. ",
        "introduction": "Gravitational microlensing has produced many interesting scientific returns concerning the Galactic structure, dark matter and variable stars (see Paczy\\'nski 1996a for a review). The current (photometric) observations of microlensing events have one notable shortcoming: from the observed time scale of microlensing event one cannot infer the lens mass and distance, since this measured quantity depends on the lens mass, distances to the lens and the source and the relative lens and source transverse motion. This degeneracy can be partially or even completely removed using some rare microlensing events (e.g., Gould 1996, 1997). A much more general way of lifting the degeneracy is conducting an accurate astrometric program together with the photometric observations ( Hosokawa et al. 1993, H${\\rm \\o}$g, Novikov \\& Polnarev 1995, Miyamoto \\& Yoshi 1995, and Walker 1995); we refer to such an astrometric experiment as astrometric microlensing. Accurate astrometry can measure the angular Einstein radius and the relative parallax between the lens and source. With these two additional constraints, the lens mass can be determined, even if we have no independent measures of the lens and source distances (Paczy\\'nski 1998; Boden, Shao, \\& Van Buren 1998). This method also opens up the possibility of measuring the mass of nearby high proper-motion stars (Miralda-Escud\\'e 1996, Paczy\\'nski 1996b). Boden et al. (1998) showed that the required astrometric accuracy can be reached by ground interferometers soon available on 8-10m class telescopes such as KECK and VLT. The Space Interferometry mission (SIM), to be launched in 2004, offers the best hope for disentangling the lens degeneracy with its global astrometric accuracy of $\\sim \\muas$. Ground and space interferometers (for a review see Shao \\& Colavita 1992) therefore offers exciting possibilities of measuring the mass of lenses, which may turn out to be brown dwarfs, planets etc.  Paczy\\'nski (1998) pointed out that, using astrometric observations, it is possible to measure the stellar radius to 1\\% accuracy for microlensed sources which show finite size effect. The photometric extended source effect has been studied by many workers (Witt \\& Mao 1994; Gould 1994; Nemiroff \\& Wickramasinghe 1994). However, up to now, most studies of astrometric microlensing adopt a point source approximation; the only exception is Walker (1995), who briefly discussed the astrometric extended source effect. Since the photometric extended source effects have already been detected for several sources (95-BLG-30, Alcock et al. 1997; 97-BLG-28, Albrow et al. 1998 and 97-BLG-56) and in light of the possibility of measuring stellar radii and other stellar parameters to excellent accuracy, we study the astrometric extended source effects in detail. In \\S 2, we first derive analytical expressions for the centroid motion for a source with uniform surface brightness. The centroid motion for sources with (realistic) limb-darkening profiles are then simplified into a one-dimensional integral, which can be computed easily numerically. In \\S 3, we discuss briefly our results. Throughout this paper, we ignore parallactic effects, since this subject has been covered in detail by other authors (e.g., Hosokawa et al. 1993, H${\\rm \\o}$g et al. 1995, Miyamoto \\& Yoshi 1995). ",
        "conclusions": "Extended source effects in photometric monitoring of microlensing events have already been observed for several events. We studied the astrometric signatures of these events in future microlensing experiments using interferometry. For microlensing events that exhibit extended source effects photometrically, the impact parameter is expected to be comparable to the source radius (e.g., Witt \\& Mao 1994); for such events, the centroid motion is significantly modified by the finite source size effect and the distinctive clover-leaf features can be easily detected with future interferometers such as SIM (cf. Figs. 1 and 2). These modifications should be taken into account when fitting the centroid motions\\footnote{The software for calculating the centroid motion of an extended source is available at \\it http://www.mpa-garching.mpg.de/\\~\\,smao/centroid}. This will add one more parameter to the fitting procedure, but this also means in principle the source size can be obtained independent of photometric observations. We conclude that using astrometric microlensing, it is possible to determine interesting stellar parameters such as stellar radius with $\\sim 1\\%$ accuracy (Paczy\\'nski 1998). We are anxiously waiting for the arrival of astrometric microlensing."
    },
    "9804/astro-ph9804273_arXiv.txt": {
        "abstract": "The gravitational microlensing light curve deviates from the ideal Paczy\\'nski form if either the lens or the source is composed of binaries: binary lens and binary source lensing events.  Currently, 6 candidate binary lens events have been reported, while the frequency of binary source events is relatively very low despite the same multiplicity of lenses and sources, and only a single candidate binary source event has been reported. To account for the rarity of binary source events, Dominik pointed out that for a typical binary source event the separation between the component source stars is very large, resulting in large difference in impact parameters between the component binary source stars.  In these cases, the light curve of the highly amplified source star is barely affected by the light from the star with low amplification, making the observed binary source light curve difficult to distinguish from that of a single source lensing event.  In this paper, we determine the fraction of events with similar source star amplifications is as much as $\\sim 8\\%$, and thus show that the very low detection rate for binary source events cannot be explained by this effect alone.  By carrying out realistic simulations of binary source events, we find that a significant fraction of binary source events are additionally missed from detection due to various other reasons. First, if the flux ratio between the component stars is very large, the light curve of the bright star is hardly affected by the light from the faint star.  Second, if the separation is too small, the binary source stars behave like a single star, making it difficult to separate the binary source event from a single event. Finally, although the probability of detecting binary source events increases as the source separation increases, still some fraction of binary source events will be missed because the light curves of these events will mimic those of single source events with longer timescales and larger values of the impact parameter. ",
        "introduction": "Experiments to detect Massive Astronomical Compact Halo Objects (MACHOs) by monitoring light variations of stars caused by gravitational microlensing have been carried out and nearly 300 events have been detected (Alcock et al.\\ 1997, Ansari et al.\\ 1996, Udalski et al.\\ 1997, Alard \\& Guibert 1997).\\markcite{alcock1997, ansari1996, udalski1997, alard1997} When both the lens and source star are approximated as point sources, the light curve of a lensing event is related to the lens-source separation in units of the Einstein ring radius $r_{\\rm E}$ by $$ A(u) = {u^2+2\\over u(u^2+4)^{1/2}};\\qquad u^2 = \\beta^2 + \\left( {t-t_0\\over t_{\\rm E}}\\right)^{2}, \\eqno(1) $$ where $\\beta$ is the impact parameter and $t_0$ is the time of maximum amplification.  The Einstein ring radius crossing time (Einstein timescale) is related to the physical parameters of the lens system by $$ t_{\\rm E} = {r_{\\rm E}\\over v};\\qquad r_{\\rm E} = \\left( {4GM\\over c^2}{D_{ol}D_{ls}\\over D_{os}}\\right)^{1/2}, \\eqno(2) $$ where $M$ is the mass of the lens, $v$ is the lens-source transverse speed, and $D_{ol}$, $D_{ls}$, and $D_{os}$ are the separations between the observer, lens, and source.  The lensing light curve deviates from the ideal form in equation (1) if either the lens or the source is composed of binaries: binary-lens (BLL) and binary-source lensing (BSL) events.  Currently, 6 candidate BLL events have been reported by various groups (MACHO LMC\\#1, Dominik \\& Hirshfeld 1994, 1996; OGLE\\#7, Udalski et al.\\ 1994; DUO\\#2, Alard, Mao, \\& Guibert \\ 1995; MACHO LMC\\#9, Bennett et al.\\ 1996; MACHO Bulge 95-12, Pratt et al.\\ 1995; MACHO Bulge 96-3, Stubbs et al.\\ 1997).\\markcite{dominik1994, 1996} In contrast, the frequency of BSL events is relatively very low despite the same multiplicity of lenses and sources, and only a single candidate BSL event has been reported (MACHO LMC 96-2, Becker et al.\\ 1997).\\markcite{becker1997} To account for the rarity of BSL events, Dominik (1998) pointed out that for a typical BSL event the separation between the component source stars is very large, resulting in large difference in impact parameters between the component binary source stars.  In these cases, the light curve of the highly amplified source star is barely affected by the light from the star with low amplification, making the observed BSL light curve difficult to distinguish from that of a single source lensing (SSL) event.  In this paper, we determine the fraction of events with similar source star amplifications is as much as $\\sim 8\\%$, and thus show that the very low detection rate for BSL events cannot be explained by this effect alone. By carrying out realistic simulations of BSL events, we find that a significant fraction of BSL events are additionally missed from detection due to various other reasons. First, if the flux ratio between the component stars is very large, the light curve of the bright star is hardly affected by the light from the faint star.  Second, if the separation is too small, the binary source stars behave like a single star, making it difficult to separate the BSL event from a SSL event. Finally, although the probability of detecting BSL events increases as the source separation increases, still some fraction of BSL events will be missed because the light curves of these events will mimic those of SSL events with longer timescales and larger values of the impact parameter. ",
        "conclusions": "The detection rate of BSL events is very low relative to the detection rate of BLL events despite similar duplicities of lenses and source stars. The scarcity of BSL event detections results from a combination of reasons. First, if the difference in impact parameters between individual components of a BSL event is very large, the light curve with lower amplification has little effect on that of the high amplification event, making the observed BSL event light curve difficult to distinguish from a SSL event one. However, we find that the fraction of BSL events with similar source star amplifications is still substantial ($\\sim 8\\%$ for $\\delta\\beta\\leq 0.1$), and thus the very low detection rate of BSL events cannot be explained by this reason alone. We find that the light curves of an important fraction of BSL events are additionally confused with SSL events due to various other reasons. These reasons include the large brightness differences between source stars, small source star separations, and the imitation by BSL events of SSL events with longer timescales and larger impact parameters."
    },
    "9804/astro-ph9804103_arXiv.txt": {
        "abstract": "We collect data at well sampled frequencies from the radio to the $\\gamma$--ray range for the following three complete--samples of blazars: the Slew Survey and the 1 Jy samples of BL Lacs and the 2 Jy sample of Flat Spectrum Radio--Loud Quasars (FSRQs). The fraction of objects detected in $\\gamma$--rays (E $\\gs 100$ MeV) is $\\sim$ 17 \\%, 26 \\% and 40 \\% in the three samples respectively. Except for the Slew Survey sample, $\\gamma$--ray detected sources do not differ either from other sources in each sample, nor from all the $\\gamma$-ray detected sources, in terms of the distributions of redshift, radio and X--ray luminosities and of the broad band spectral indices (radio to optical and radio to X--ray).  We compute average spectral energy distributions (SEDs) from radio to $\\gamma$--rays for each complete sample and for groups of blazars binned according to radio luminosity, irrespective of the original classification as BL Lac or FSRQ.  The resulting SEDs show a remarkable continuity in that: i) the first peak occurs in different frequency ranges for different samples/ luminosity classes, with most luminous sources peaking at lower frequencies; ii) the peak frequency of the $\\gamma$--ray component correlates with the peak frequency of the lower energy one; iii) the luminosity ratio between the high and low frequency components increases with bolometric luminosity.  The continuity of properties among different classes of sources and the systematic trends of the SEDs as a function of luminosity favor a unified view of the blazar phenomenon: a single parameter, related to luminosity, seems to govern the physical properties and radiation mechanisms in the relativistic jets present in BL Lac objects as well as in FSRQ. The general implications of this unified scheme are discussed while a detailed theoretical analysis, based on fitting continuum models to the individual spectra of most $\\gamma$-ray blazars, is presented in a separate paper (Ghisellini et al. 1998). ",
        "introduction": "The discovery of BL Lac objects and the paradoxes associated with their violent variability led to a major step forward in the theory of Active Galactic Nuclei (AGN), that is to the concept of relativistic jets. Flat spectrum, radio--loud quasars (Angel \\& Stockman 1980) share basically all of the properties of BL Lac objects related to the presence of a strong non--thermal broad band continuum, except for the absence of broad emission lines. Hence the common designation of blazars proposed by Ed Spiegel in 1978.  It was initially supposed that BL Lacs represented the most extreme version of FSRQs, i.e. those with the most highly boosted continuum. Instead, it has been recognized later (e.g. Ghisellini, Madau \\& Persic 1987; Padovani 1992a; Ghisellini et al. 1993) that the {\\it amount of relativistic beaming} and the {\\it intrinsic} power in the lines are lower in BL Lacs than in FSRQs, implying some intrinsic difference between the two classes. Differences are also found in the extended radio emission and jet structure (e.g. Padovani 1992a; Gabuzda et al. 1992). Nevertheless the continuity of several observational properties including the luminosity functions (Maraschi \\& Rovetti 1994), the radio to X--ray SEDs (Sambruna et al. 1996) and the luminosity of the lines (Scarpa \\& Falomo 1997) suggests that blazars can still be considered as a single family, where the physical processes are essentially similar allowing for some scaling factor(s).  The identification of these scaling factors would represent a substantial progress in the understanding of the blazar phenomenon.  A special class of BL Lacs was found from identification of  X--ray sources. X--ray selected BL Lacs (XBL) differ from the classical radio--selected BL Lacs (RBL) in a lesser degree of ``activity'' (including polarization), in the radio to optical emission and in the relative intensity of their X--ray and radio emission. This led to the suggestion that the X--ray radiation was less beamed than the radio one and that XBLs were observed at larger inclination to the jet axis (e.g. Urry \\& Padovani 1995 for a review).  Giommi \\& Padovani (1994) quantified the differences in SEDs between XBLs and RBLs, and Padovani \\& Giommi (1995) introduced the distinction between 'High--energy peak BL Lacs' (HBL) and 'Low--energy peak BL Lacs' (LBL), for objects which emit most of their synchrotron power at high (UV--soft--X) or low (far--IR, near--IR) frequencies respectively. Quantitatively a distinction can be done on the basis of the ratio between radio and X-ray fluxes (see also \\S3.2.2). We will use the broad band spectral index $\\alpha_{\\rm RX}$\\footnote{Hereinafter we define spectral indices as F$_\\nu \\propto \\nu^{-\\alpha}$. In broad band indices radio, optical and X--ray fluxes are taken at 5 GHz, 5500 \\AA\\ and 1 keV, respectively.} and call HBL and LBL objects having  $\\alpha_{\\rm RX} \\ls 0.75$, $\\gs 0.75$, respectively. Giommi and Padovani also proposed that HBL represent a small fraction of the BL Lac population and are numerous in X--ray surveys only due to selection effects. An alternative hypothesis (Fossati et al. 1997) relates the spectral properties to the source luminosity in such a way that low luminosity objects (with high space density) are HBLs while high luminosity objects (with low space density) are LBLs.  We will include here X--ray selected BL Lacs together with ``classical'' BL Lacs in the blazar family, again assuming that the basic physical processes by which the continuum is produced are common to the whole family.  The detection by  EGRET, on board the Compton Gamma--Ray Observatory (CGRO), of many  blazars at $\\gamma$--ray energies (E $\\gs 30$ MeV) revealed that a substantial fraction and in some cases the bulk of their power is emitted in this very high energy band. The $\\gamma$--ray emission  is therefore of fundamental importance in the SED of blazars.  From the theoretical point of view the radio to UV continuum is universally attributed to synchrotron emission from a relativistic jet, while a flat inverse Compton component due to upscattering of the low energy photons is expected to emerge at high energies as originally discussed in Jones, O'Dell \\& Stein (1974). The latter process is therefore a plausible candidate to explain the $\\gamma$--ray emission. The soft photons to be upscattered could be either the synchrotron photons themselves (synchrotron self--Compton process, SSC, e.g. Maraschi, Ghisellini \\& Celotti 1992; Bloom \\& Marscher 1993) or photons produced by the disk and/or scattered /reprocessed in the broad line region (Blandford 1993; Dermer \\& Schlickeiser 1993; Sikora, Begelman \\& Rees 1994; Ghisellini \\& Madau 1996). Understanding whether/how the $\\gamma$--ray properties differ among subclasses is essential to assess the role of different mechanisms and to verify whether the idea of blazars as a unitary class can be maintained.  Here we study the systematics of the SEDs of blazars using data from the radio to the $\\gamma$-ray band. We confirm and extend previous results of Maraschi et al. (1995) and Sambruna et al. (1996) by: i) {\\it extending the SED to the $\\gamma$-ray range}; ii) using a much larger complete sample of FSRQ; iii) using the richer and brighter sample of X--ray selected BL Lacs recently derived from the Slew Survey. We use the available $\\gamma$--ray data for each sample but also indirect information derived from the $\\gamma$--ray detected (not complete) sample discussed by Comastri et al. (1997). Since we find that the continuity hypothesis among blazars holds we also consider a merged ``global'' sample subdivided in luminosity bins irrespective of the original classification of the objects.  The structure of the paper is as follows. In section 2 we describe how the data for the SEDs were collected and treated for each sample.  The $\\gamma$--ray properties of different samples are also discussed. In section 3 we build average SEDs for the three sub--samples and for the global sample subdivided according to luminosity and present our results. These are discussed in section 4, while our conclusions are drawn in section 5. ",
        "conclusions": "The main conclusion of this work is that despite the differences in the continuum shapes of different sub--classes of blazars, a unitary scheme is possible, whereby Blazar continua can be described by a family of analytic curves with the source luminosity as the fundamental parameter.  The \"scheme\" (admittedly empirical) determines both the frequency and luminosity of the peaks in the synchrotron and inverse Compton power distributions (and therefore also the  $\\gamma$--ray luminosity in the EGRET range) starting from the radio luminosity only. The main suggested trend is that with increasing luminosity both the synchrotron peak and the inverse Compton peak move to lower frequencies and that the latter becomes energetically more dominant. The scheme is testable, for instance it predicts that sources emitting strongly in the TeV band have relatively low intrinsic luminosity.  The \"spectral sequence\" finds a plausible interpretation in the framework of relativistic jet models radiating via the synchrotron and inverse Compton processes if the physical parameters (magnetic field and/or critical energy of the radiating electrons) vary with luminosity or if  photons outside the jet become increasingly important as seed photons for the inverse Compton process in sources of larger luminosity. The latter alternative is supported at least qualitatively by the increasing dominance of emission lines in higher luminosity objects.   The proposed scenario, in which the intrinsic jet power regulates, in a continuous sequence, the observational properties from the weaker HBL, through LBL, to the most powerful FSRQs, also fits in very nicely with the unification of FR~I and FR~II type radio galaxies as proposed by Bicknell (1995). After a long debate the prevailing view is that FR~I and FR~II radio galaxies both contain relativistic jets which can be decelerated giving rise to the FR~I morphology depending on the kinetic power in the jet and the pressure of the ambient medium.  The whole radio--loud AGN population could be unified in a two parameter space one being the intrinsic jet power, the other the viewing angle. An interesting point for future discussion is whether a third parameter associated with the luminosity of an accretion disk is necessary or is already implicitly and uniquely linked to the jet power."
    },
    "9804/astro-ph9804191_arXiv.txt": {
        "abstract": "We have reconstructed the spectrum of the afterglow of GRB 970508, on May 21.0 UT (12.1 days after the GRB), on the basis of observations spanning the X-ray to radio range.  The low-frequency power law index of the spectrum, $\\alpha=0.44\\pm0.07$ ($F_\\nu\\propto\\nu^\\alpha$), is in agreement with the expected value $\\alpha=1/3$ for optically thin synchrotron radiation. The 1.4 GHz emission is self-absorbed. We infer constraints on the break frequecies $\\nu_{\\rm c}$ and $\\nu_{\\rm m}$ on May 21.0 UT from a spectral transition from $F_{\\nu} \\sim \\nu^{-0.6}$ to $F_{\\nu} \\sim \\nu^{-1.1}$ in the optical passband around 1.4 days. A model of an adiabatically expanding blast wave, emitting synchrotron radiation, in which a significant fraction of the electrons cool rapidly provides a successful and consistent description of the afterglow observations over nine decades in frequency, ranging in time from trigger until several months later. ",
        "introduction": "GRB970508 was a moderately bright $\\gamma$-ray burst (Costa et al. 1997; Kouveliotou et al. 1997). It was detected on May 8.904 UT with the Gamma-Ray Burst Monitor (GRBM; Frontera et al. 1991), and with the Wide Field Cameras (WFCs; Jager et al. 1995) on board the Italian-Dutch X-ray observatory BeppoSAX (Piro et al. 1995). Optical observations of the WFC error box (Heise et al. 1997), made on May 9 and 10, revealed a variable object at RA = $06^{\\rm h}53^{\\rm m}49\\fs2$, Dec = +79\\degr16\\arcmin19\\arcsec (J2000), which showed an increase by $\\sim$1 mag in the V band (Bond 1997), and whose spectrum indicated that the distance to the GRB source corresponds to a redshift of at least 0.83 (Metzger et al. 1997). The BeppoSAX Narrow Field Instruments (NFIs) revealed an X-ray transient (Piro et al. 1997a) whose position is consistent with that of the optical variable, and Frail et al. (1997a,b) found the first GRB radio afterglow for GRB 970508; the radio source position coincides with that of the optical source (Bond 1997).  The global properties of the X-ray, optical and radio afterglows of GRBs have been successfully described by relativistic blast wave models (e.g. Wijers, Rees and M\\'esz\\'aros 1997). In this paper we show that the  detailed properties of the spectrum of the afterglow of GRB 970508, and its evolution, are well described by a blast wave model in which significant cooling of the electrons plays an important role (Sari, Piran and Narayan 1998). This model predicts the occurrence of an extra break at $\\nu_{\\rm c}$ in the spectrum, corresponding to the Lorentz factor $\\gamma_{\\rm c}$ above which the electron cooling time is shorter than the expansion time of the blast wave. ",
        "conclusions": "We have found that on May 21.0 UT the afterglow spectrum of GRB 970508 contains all the characteristic parts of a synchrotron spectrum. We infer that the dynamical evolution of the remnant is adiabatic (in agreement with Waxman et al. 1998), but that a significant fraction of the electrons have synchrotron cooling times shorter than the remnant's expansion time. In accordance with the theory outlined by Sari et al. (1998)  this produces an additional break in the spectrum. The break frequencies $\\nu_{\\rm m}$ and $\\nu_{\\rm c}$, corresponding to the minimum electron energy and the energy above which electrons cool rapidly, respectively, are equal at a time $t_0$ which we find is 500 s. This synchrotron spectrum and adiabatic dynamics can explain most of the afterglow behaviour over 9 decades in frequency spanning the X rays to radio waves, from the first optical to the late radio data. Some aspects of the afterglow are not explained by this model: the flux at $\\nu_{\\rm m}$ should be constant, but is seen to decrease from 12 days when it lies at mm wavelengths to 55 days when it lies in the radio; the self-absorption frequency $\\nu_{\\rm a}$ is predicted to be time-independent, but is seen to decrease with time and the early rise of the radio fluxes is slower than expected (see Paper I) and finally the decay after maximum at mm wavelengths is perhaps somewhat faster than expected (see paper I for a discussion on the radio light curves).  In summary, we find that an adiabatically expanding blast wave emiting synchrotron radiation, with a significant fraction of rapidly cooling electrons, describes most of the afterglow data on GRB 970508 very well."
    },
    "9804/astro-ph9804158_arXiv.txt": {
        "abstract": "s{We use current cosmic microwave background data to constrain cosmological parameters.  The results are presented as confidence regions in the ($\\Omega$--$H_o$) -- plane for both open and flat geometries, and are based on a $\\chi^2$ minimization to the power spectrum. Although such an approach does not strictly apply to power spectrum estimates, the results should nonetheless be indicative. We find that purely open, low--$\\Omega$ models are strongly disfavored, due to the position of the ``Doppler Peak'' at high $l$; flat models are less strongly constrained.  The parameter space explored is the largest to date, covering $\\Omega$, $H_o$, $n$, $Q$ and $\\Omega_B h^2$.} ",
        "introduction": "It is already possible, with existing data on the fluctuations of the cosmic microwave background (CMB), to constrain some cosmological parameters (Rocha \\& Hancock 1996; Bond \\& Jaffe 1996; Lineweaver et al. 1997).  Of course, the actual data is far from providing the kind of constraints eagerly awaited from the next generation of experiments (Jungman et al. 1996; Bond et al. 1997), but there are nevertheless more than 10 different experiments detecting fluctuations over a large range of angular scales. To place our present work in context, consider the general problem of analyzing the results (say, a map) of a CMB experiment.  If the true sky fluctuations and, as is fortunately often the case, the instrumental noise are gaussian, then the probability of obtaining any particular set of $N$ pixel values is a multivariate gaussian in the pixel vector $\\vec{d}$ (a one--dimensional listing of the possible pixel values): \\begin{displaymath} {\\cal L}(\\vec{p}) = {\\cal P}(\\vec{d}/\\vec{p}) = \\frac{1}{(2\\pi)^{N/2} |C|^{1/2}} e^{-(\\vec{d})^t \\cdot C^{-1} \\cdot (\\vec{d})/2} \\end{displaymath} The cosmological model is completely embodied (as well as the noise) in the covariance matrix of the pixels, $C \\equiv C^T(\\vec{p}) + C(noise)$, where we denote the underlying model parameters by the vector $\\vec{p}$. Thus, we have a likelihood function for $\\vec{p}$, given a particular $\\vec{d}$, i.e., a map of the sky.  It is important to note that even if the pixels are gaussian, the likelihood function ${\\cal L}(\\vec{p})$ is NOT so, in general, because the parameters enter into the {\\em covariance} matrix and not as linear combinations of the pixel values themselves.  Maximization of this likelihood function is one way to estimate cosmological parameters.  Another approach is to work in Fourier space with estimates of the power spectrum, i.e., the $C_l$. These are nothing more than the variances of the individual spherical harmonic coefficients of the temperature fluctuations.  We are really doing the same thing as before, just working with the covariances of Fourier coefficients instead of the measured pixel values.  A practical reason for using this method is that band--power estimates and their uncertainties are readily available in the literature, this being the most common way of concisely reporting the results of an experiment.  In the present contribution, we assemble various band--power estimates (e.g., $C_l$ intergrated over experimental window functions) and constrain certain cosmological parameters by fitting to these power spectrum data points.  We must mention several caveats applying to the results given here: Firstly, the data set is certainly not homogeneous, if for no other reason than for the different ways that the so--called cosmic variance has been estimated and included in the quoted error bars. Secondly, we have not yet fully accounted for the experimental window functions.  A third important point is that the $C_l$ are NOT gaussian random variables, unlike the individual pixel values (or Fourier coefficients). This is clear -- the $C_l$ represent the variance of gaussian random variables, and thus are themselves distributed according to a $\\chi^2$ distribution. We have, for the work presented in this contribution, effectively assumed gaussianity in the sense that we apply a $\\chi^2$ minimization for the fitting.  In principle, this could lead to biased best--estimates and incorrect confidence intervals.  We are currently working on improving these aspects of our calculations. The results presented here should therefore be taken as indicative, but perhaps relatively good indications, all the same, given the current status of affairs (Jaffe \\& Bond 1997).\\\\  \\begin{tabular}{|c||c|c||c|c|} \\multicolumn{5}{c}{\\bf\\large \\hspace*{1cm} Open Models\\hspace*{4.5cm} Flat Models} \\\\ \\hline\\hline parameter & range & step & range & step \\\\ \\hline\\hline $\\Omega$ \t& 0 -- 0.95\t \t& 0.05\t& 0 -- 1\t& 0.05 \\\\ $H_o$\t\t& 15 -- 100 km/s/Mpc \t& 5 km/s/Mpc & 15 -- 100 km/s/Mpc \t& 5 km/s/Mpc \\\\ $n$\t\t& 0.5 -- 1.5\t\t& 0.03\t& 0.55 -- 1.45\t& 0.09 \\\\ $Q$\t\t& 14 -- 20 $\\mu$K\t& 1 $\\mu$K & 14 -- 20 $\\mu$K \t& 1 $\\mu$K\\ \\\\ $\\Omega_B h^2$\t&    0.015\t\t& ---- & 0.006 -- 0.030\t& 0.002  \\\\ \\hline\\hline \\end{tabular} ",
        "conclusions": "It is a little early to draw detailed conclusions from the CMB concerning the cosmological parameters, but it is worth noting that some conclusions are already possible with existing data; and the next round of data releases will not be long in coming. Despite the various caveats of the method used here, it seems that one result should remain rather robust - that low--$\\Omega$, purely open models are strongly disfavored, for the clear and simple reason that the ``Doppler Peak'' is just not in the correct place due to the large angular distance to the surface of last scattering in such models."
    },
    "9804/astro-ph9804228_arXiv.txt": {
        "abstract": "We report results from a long-term monitoring data set obtained with the Rossi X-ray Timing Explorer (RXTE) from the bright globular cluster binary source 4U 1820-30.  During this monitoring, the source intensity varied from a flux of $2.9\\times 10^{-9} ergs\\  cm^{-2} s^{-1}$ to $8.5\\times 10^{-9} ergs\\  cm^{-2} s^{-1}$ in the 2-10 keV band, which correspond to, respectively, a count rate of 1,600 cps and 4,500 cps in the nominal 2-60 keV RXTE-PCA band. Its energy spectral state was observed to encompass both the banana branches and the island state. Kilohertz QPOs were observed to exist in the upper-banana and lower-banana branches, as well as in the island state.  In particular, we report that the kHz QPOs were present only when the soft-color (defined as the flux ratio of the 3.9-5.0 keV band to the 1.7-3.9 keV band) was below 0.78 without regard to the overall count rate and without regard to whether the source is in any one of the three branches or states.  Their centroid frequency was well correlated with the overall count rate below 2,500 cps. When the count rate was above 2,500 cps, the QPO frequency was consistent with a constant independent of the overall count rate. We take this as further evidence that the inner edge of the accretion had reached the marginally stable orbit, and therefore the constant 1060$\\pm 20$ Hz is the marginally stable orbit frequency. ",
        "introduction": "Low mass X-ray binaries are enigmatic objects. They are the brightest X-ray sources in the sky. There is little  doubt that they harbor fast-spinning neutron stars, yet all efforts to date of searching for coherent pulsations in their persistent flux have resulted in negative results (see, e.g. \\cite{Vaughan_etal_1994}).  Those searches, however, have uncovered quasi-periodic flux oscillations which are believed to be indirectly related to the neutron star spin (\\cite{mbfm,vanderKlis_1985}).  Most recently, since their discovery with RXTE in 4U 1728-34 (\\cite{4u1728_discovery}) and Sco X-1 (\\cite{scox1_discovery}), the highest frequencies observed from any cosmic sources, the so-called kilohertz quasi-periodic oscillations, have been observed from nearly every LMXB.  For a comprehensive review of the literature on this subject, see van der Klis (1997).  The  bright globular cluster binary 4U 1820-30 is one of the most studied low mass X-ray binaries because of its membership in a globular cluster, its Type-I X-ray bursts (\\cite{discovery_of_bursts}), its variability on all time scales, its short orbital period of only 685 s, and most recently, the kilohertz quasi-periodic oscillations in the persistent flux (\\cite{4u1820_discovery}).  Since it is in the globular cluster NGC 6624, its distance of $6.4\\pm 0.6$ kpc is probably more accurately known to us than any other LMXB. Its short orbital period--- the shortest known of any cosmic binary---strongly suggests that its mass-donating star is most likely a helium dwarf with 0.06 to 0.08 solar masses (\\cite{Rappaport_etal_1987}).  The characteristics of the kilohertz QPOs discovered in the persistent flux (\\cite{4u1820_discovery}) were rather similar to those of other bursters. Their centroid frequency was quite tightly correlated with the 2-60 keV count rate.  Their fractional rms amplitudes were of the order of 3\\%. The frequency difference of the twin QPO peaks was consistent with a constant $275\\pm 8$ Hz.  Their fractional rms amplitudes increase dramatically with energy; the 8-20 keV amplitudes are 1.8-2.9 times greater than those measured in the 1-8 keV band.  In this paper we report on data sets that sample the large count rate and energy spectral span of this source over a calendar period of 6 months starting in February and ending in September of 1997.  In the discussion section, we place our results in the overall context of kHz QPO phenomenologies and discuss the most important features of our results. ",
        "conclusions": "We have observed kHz QPOs from 4U 1820-30 on both the banana branches and the island state. This is the first source from which kHz QPOs have been observed in all branches or states (see \\cite{4u1636_2qpos} and \\cite{4u1735_discovery} for discussions).  The existence of kHz QPOs in the flux is strongly correlated with the source position on the color-color diagram. We have examined the correlation between the source states as defined by Hasinger and van der Klis (1989) and the existence of kHz QPOs and found that there does not seem to be a one to one correspondence between the two. For example, we found kHz QPOs in all three branches/state in some segments of the data and found no kHz QPOs in any of the three branches/state in other segments of the data. Therefore we think that over long terms color-color daigram alone may be a better predictor or indicator for the existence of kHz QPOs.  Figure~\\ref{f_vs_cps} demonstrates the well-defined correlation between the QPO centroid frequency and the overall count rate. It is similar to nearly all other sources for which this correlation is known. Although the data analyzed here covers a large span of calendar time, there were no large shifts in this correlation, in contrast to the cases of 4U 0614+091 (\\cite{ford_etal_4u0614}), 4U 1608-52 (\\cite{yu_etal_1608}), and Aql X-1 (\\cite{aqlx1_discovery}). A very important feature of Figure~\\ref{f_vs_cps}, compared to the frequency vs. count rate correlations of other sources, is that the QPO frequency is independent of the count rate in the range of 2,600 to 3,200 cps.  This is a feature that has been proposed (\\cite{spm}) to strongly support the hypothesis (\\cite{kaaret_etal_1997,NS_mass_estimate_1}) that the highest observed QPO frequencies are the marginally stable orbit frequencies near the neutron stars in the low mass X-ray binaries.  Figure~\\ref{f_vs_cps} strongly suggests that indeed we have observed the highest QPO frequency from 4U 1820-30. If this frequency is the marginally stable orbit frequency, then neutron star in 4U 1820-30 would have a mass in the vicinity of 2.2$M_\\odot$. Given the importance of this conclusion, more corroborating evidence from other sources is needed.  Figure~\\ref{qpo_vs_cps}b tests the constancy of the frequency difference over a much larger span of count rate than previously possible (\\cite{4u1820_discovery}). With the statistical accuracy we can measure, the difference is a constant 275$\\pm$8 Hz.  This is similar to the cases of 4U 1728-34 and 4U 0614+091, where the differences are consistent with being constants 363 and 327 Hz (\\cite{4u1728_discovery,4u0614_discovery}), respectively; in contrast to the cases of Sco X-1 and 4U 1608-52 whose frequency differences have been observed to change by over 100 Hz (\\cite{scox1_separation,4u1608_2nd_peak}).  The fractional rms amplitude as a function count rate, Figure~\\ref{qpo_vs_cps}c, indicates that, on average, when the count rate decreases by a factor of 2 from 3,200 cps to 1,600 cps, the rms amplitude increases by a factor of 2 from about 3\\% to about 6\\%. This indicates that the flux pulsed at the QPO frequency has a constant amplitude. It is only the un-pulsed flux that has changed by a factor of 2. This point can be potentially constraining as it requires that the overall X-ray luminosity of the source change by as much as a factor of 2, but the flux pulsed at the kHz frequencies remain more or less the same.  As the source count rate decreses, the upper frequency QPO peak becomes broader and its amplitude larger.  The lower frequency QPO would be sandwiched in between the low frequency noise which is below 100 Hz and the upper QPO if it were below the upper frequency by 275 Hz.  We have no positive detection of it and set a 90\\% confidence level upper limit on the fractional rms amplitude of 2\\%.  Finally we comment on the transition of the source from the banana branch to the island state. As indicated in the original classification paper (\\cite{hk89}), our observation strongly suggests that the transition from the banana branch to the island state is most likely a continuous one (see the last paragraph in the last section), not an abrupt change. As the source moves closer to the island state, i.e., when the count rate decreses, the upper QPO peak becomes broader and its overall fractional rms amplitude becomes higher.  Its centroid frequency, to the extent that we can tell from our data sets, follows continuously along the expected curve extroplated from the banana branch."
    },
    "9804/astro-ph9804252_arXiv.txt": {
        "abstract": "We present  a study of  recent  solar neutrino  data using a  Bayesian method.    Assuming   that   only   $\\nu_e$   are   observed   in  the Super-Kamiokande  experiment our results  show a  marked supression of the survival probability at about 1~MeV, in  good agreement with $\\chi ^2$-based   analyses.   When      the detection  of    $\\nu_{\\mu}$  by Super-Kamiokande is    taken   into  account,   assuming $\\nu_e$    to $\\nu_{\\mu}$ oscillations, we find the  largest suppression in survival probability at about 8.5~MeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804064_arXiv.txt": {
        "abstract": "We report redshifts, magnitudes and morphological classifications for 5369 galaxies with $m_B \\leq 15.5$ and 57 galaxies fainter  than this limit, in two regions covering a total of 1.70 steradians in the southern celestial hemisphere. The galaxy catalog is drawn primarily from the list of non-stellar objects identified  in the Guide Star Catalog (Lasker et al. 1990, AJ 99, 2019; hereafter GSC). The galaxies have positions accurate to $\\sim 1^{\\prime\\prime}$ and magnitudes with an rms scatter of $\\sim 0.3^m$. We compute magnitudes ($m_{SSRS2}$) from the relation  between instrumental GSC magnitudes and the photometry by Lauberts \\& Valentijn (1989). From a comparison with CCD photometry, we find that our system is homogeneous across the sky and corresponds to magnitudes measured at the isophotal level $\\sim$ 26 mag arcsec$^{-2}$. The precision of the radial velocities is of $\\sim$ 40 \\kms\\ and the redshift survey is more than  99\\% complete  to the $m_{SSRS2}$ = 15.5 magnitude limit.  This sample is in the direction opposite to the CfA2; in combination the two surveys provide an important database for studies of the properties of galaxies and their large-scale distribution in the nearby Universe. ",
        "introduction": "During the past two decades, complete, wide-angle redshift surveys have provided a fundamental basis for exploring the large-scale distribution of galaxies. Starting with the  CfA1 Redshift Survey (Davis \\etal 1982), the first to reach beyond the Local Supercluster, these surveys have revolutionized our view of the spatial distribution of galaxies. The CfA1 redshift survey provided the first three-dimensional glimpse of the nearby universe, showing that within 70$h^{-1}$ Mpc ($H_0$ =100\\h1 km/s/Mpc) the galaxy distribution is far from homogeneous with large regions devoid of bright objects. The first complete slice of the CfA2 Redshift Survey (de Lapparent et al. 1986) soon produced striking results with the discovery of the Great Wall, a spectacular example of a well-defined, thin, two-dimensional structure perpendicular to the line-of-sight and  outlining several large voids (de Lapparent et al. 1986). Later the diameter-limited Southern Sky Redshift Survey (SSRS, da Costa et al. 1988), not only confirmed the existence of large voids, but also detected other thin two-dimensional coherent structures, like the Southern wall.  These surveys revealed the large scale of structure in the universe.  Surveys of both hemispheres have continued: the CfA2 covers a portion of both galactic caps in the northern celestial hemisphere (Geller \\& Huchra 1989) and the SSRS2 (\\eg da Costa et al. 1994$a$) covers a portion of both caps in the south. In contrast to the diameter-limited SSRS, the deeper SSRS2 is a magnitude limited survey based on a selection of objects from the Hubble Space Telescope Guide Star Catalog (Lasker \\etal 1990; GSC hereafter).  The CfA2 and SSRS2 surveys now cover more than 30\\% of the sky. Together they contain more than 20,000 galaxies and provide a panoramic view of the three-dimensional galaxy distribution  (\\eg da Costa \\etal 1994$a$). The redshift survey maps show that the galaxy distribution consists of a complex network of voids, typically 5000 \\kms in size, surrounded by thin walls where most galaxies reside (Geller \\& Huchra 1989).  A quantitative analysis of  the SSRS2 (El-ad \\etal 1996) confirms this picture, finding the typical size of voids to be $\\sim$ 40\\h1 Mpc. Contrary to some earlier  expectations, large voids and walls are not rare events; they are  present in every direction surveyed (\\eg CfA2; Las Campanas Redshift Survey, Schectman et al. 1996).  Both separately and combined, the CfA2 and SSRS2 have provided a basis for determination of the local luminosity function (Marzke \\etal 1994$b$, da Costa \\etal 1994$a$) and its dependence on morphology (Marzke \\etal 1994$a$, Marzke \\etal 1998) and color (Marzke \\& da Costa 1997); the power-spectrum of the galaxy distribution (Park \\etal 1994, da Costa \\etal 1994$b$); the mean pairwise velocity dispersion (Marzke \\etal 1995); the dependence of clustering properties on luminosity (Benoist \\etal 1996), morphology and color (Willmer et al. 1998) and the properties  of loose and compact groups (Barton \\etal 1996, Ramella \\etal 1997, Ramella \\etal 1998).  Even though the currently planned deeper surveys will dramatically extend the depth of the sampled volume of the Universe, the combined CfA2 and SSRS2, which probe the distribution of $L_*$  galaxies out to $\\sim$ 100 \\h1 Mpc, will continue to be a valuable database. For instance, the large number of galaxies, the complete sampling and the large angular coverage will allow, in conjunction with independent distance measurements, a more detailed study of the relation between individual galaxies and large-scale structure.   Here we provide the SSRS2 catalog. In Section 2 we discuss the assembly of the sample of galaxies from the GSC, the magnitude system, and  the morphological classification of galaxies. In Section 3 we describe the observations. The catalog itself is in Section 4. A brief summary follows in Section 5. ",
        "conclusions": "The SSRS2 catalog, described in this paper, includes 5426 galaxies in the southern hemisphere.  For essentially all of these galaxies we present  radial velocities, good astrometric positions and objectively measured  magnitudes. The catalog also includes galaxy morphologies compiled from  a variety of sources.  The magnitude limited sample designed for  statistical analyses contains 5369 galaxies brighter than $m_{SSRS2}$ = 15.5, distributed within a solid angle of 1.7 steradians,  split between two regions in the north and south galactic caps. The  magnitude system is homogeneous and equivalent to isophotal, measured at  26 mag arcsec$^{-2}$. Its zero-point offset relative to the Zwicky  magnitudes, used in the CfA surveys, is estimated to be about 0.1$^m$.  Since the loss of objects at faint magnitudes, caused by misclassifications of galaxies as stars in the source catalog (GSC), is small, we conclude that the magnitude limited catalog is essentially complete to $m_{SSRS2}$ = 15.5.  The redshift completeness exceeds 99\\%.  The SSRS2 galaxy catalog is the basis for a variety of studies of large-scale structure either by itself or in combination with the CfA2 to  probe larger scales.  Together these surveys are the largest database of  redshifts of the Local Universe.  The authors thank the time allocation committees of all observatories, for granting this project a generous amount of dark time over a number of years. More recent observations were also made possible thanks to the ESO/Observat\\'orio Nacional agreement for shared operation of the 1,52 m ESO telescope at La Silla. The authors thank S. Paoloantonio, C. Valotto and J.H. Calder\\'on for their help in the observations at CASLEO, and M.A. Nunes and D. Nascimento for their help maintaining the ON Reticon in working order. We also thank Will Sutherland, Steve Maddox and Somak Raychaudhuri for providing data from APM surveys.  At the CfA, we thank Barbara Elwell for her help in the visual inspection stage of the catalog construction, Susan Tokarz for spectroscopic data reduction, and Cathy Clemens for her help in the use of the ZCAT database.  LNdC appreciates the hospitality of CfA and IAP where part of this work was carried out, and acknowledges support from a CNPq fellowship, the Smithsonian Scholarly Program and the John Simon Guggenheim Foundation. This project has received support from CNPq-NSF (Brazil-USA) and CNPq-CONICET (Brazil-Argentina) bilateral agreements. Partial support has also been provided through CNPq grants 301364/86-9, 453488/96-0 (CNAW), 301373/86-8 (PSP). Additional funds were provided by the Centro Latino-Americano de F\\'\\i sica (PSP), FAPERJ (OLC), NSF AST-9023178 (CNAW) and the ESO Visitor Programme (CNAW).  The authors acknowledge use of the CCD and data acquisition system supported under U.S. National Science Foundation grant AST-90-15827 to R.M. Rich.  This work has made extensive use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration and of the NASA Astrophysics Data System.  This work has also used images from the Digitized Sky Survey,  produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope.  \\clearpage"
    },
    "9804/astro-ph9804314_arXiv.txt": {
        "abstract": "A total of 235 Active Galactic Nuclei (AGN) from two different soft X-ray surveys (the \\ROSAT\\ Deep Survey --DRS-- and the \\ROSAT\\ International X-ray Optical Survey --RIXOS--) with redshifts between 0 and 3.5 are used to study the clustering of X-ray selected AGN and its evolution.  A 2~sigma significant detection of clustering of such objects is found on scales $<40-80\\hmpc$ in the RIXOS sample, while no clustering is detected on any scales in the DRS sample. Assuming a single power law model for the spatial correlation function (SCF), quantitative limits on the AGN clustering have been obtained: a comoving correlation length $1.5\\approxlt r_0 \\approxlt 3.3\\hmpc$ is implied for comoving evolution, while $1.9\\approxlt r_0 \\approxlt 4.8$ for stable clustering and $2.2\\approxlt r_0 \\approxlt 5.5$ for linear evolution.  These values are consistent with the correlation lengths and evolutions obtained for galaxy samples, but imply smaller amplitude or faster evolution than recent UV and optically selected AGN samples.  We also constrain the ratio of bias parameters between X-ray selected AGN and $IRAS$ galaxies to be $\\approxlt 1.7$ on scales $\\approxlt 10\\hmpc$, a significantly smaller value than is inferred from local large-scale dynamical studies. ",
        "introduction": "Since the launch of \\ROSAT\\ in 1991, a large number of surveys of soft X--ray selected sources have been undertaken using the Position Sensitive Proportional Counter (PSPC), with different sky coverages, depths and identification completenesses.  They range from the \\ROSAT\\ All Sky Survey (RASS, Voges 1992) that covers the whole sky to a fairly bright flux limit, passing through the relatively large solid angle RIXOS (20 $\\deg^2$ Mason et al., in preparation), to the deeper \\ROSAT\\ Deep Survey (henceforth DRS, Shanks et al. 1991), Cambridge Cambridge \\ROSAT\\ Serendipitous Survey (Boyle et al. 1995) or the UK Medium Survey (Carballo et al. 1995). The deepest and narrowest surveys, pushing to the limit the capabilities of the PSPC are the UK Deep Survey (Branduardi--Raymont et al. 1994, McHardy et al. 1998) and the Deep Survey in the Lockman Hole (Hasinger et al. 1993, Hasinger et al. 1998, Schmidt et al. 1998).  Together, these surveys have given considerable insight into the evolution of the AGN X--ray luminosity function (Boyle et al. 1994, Page et al.  1996, Jones et al. 1997), their X--ray spectra (Bade et al. 1995, Carballo et al. 1995, Almaini et al. 1996, Romero--Colmenero et al. 1996,  Ciliegi et al. 1997, Mittaz et al. 1998) and their relation to Narrow Emission Line Galaxies (NELGs, Boyle et al. 1995, Griffiths et al. 1996, Page et al. 1997).  However, so far the only direct study of the clustering properties of X--ray selected AGN is that by Boyle \\& Mo (1993, henceforth BM). They studied the local ($z<0.2$) AGN in the Extended Einstein Medium Sensitivity Survey (EMSS, Stocke et al. 1991) and found no signal above 1 sigma, consistent with the clustering properties of the UV and optically selected AGN.  In this work we present for the first time a direct study of the clustering of soft X--ray selected AGN from \\ROSAT, using RIXOS and the \\ROSAT\\ Deep Survey (positions and redshifts for the AGN identified in the five DRS fields have been kindly provided by O. Almaini, prior to publication).  The samples are introduced in Section \\ref{data}. We discuss the evidence for clustering in Section \\ref{clust}, using different model independent methods. Quantitative measurements on the correlation strength and evolution are then obtained in Section \\ref{calscf}. These results are presented in Section \\ref{limits}, and discussed in Section \\ref{discuss}. We summarize our conclusions in Section \\ref{summa}.  We have used $H_0=100\\,h$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ (unless otherwise stated) throughout this paper. ",
        "conclusions": "\\label{summa}     The spatial correlation function of soft X--ray selected AGN has been studied here using data from two different X--ray surveys: the \\ROSAT{ } Deep Survey (DRS) and RIXOS.  Some indication of clustering at the 2~sigma level in RIXOS in the $z=0.5-1$ redshift range has been found using a variant of the counts--in--cells method.  A more powerful test has also been performed, by comparing the number of pairs of sources in each sample with the number of pairs expected from a uniform distribution of sources, finding a $\\sim$2~sigma significant clustering signal in the whole RIXOS sample ($\\langle z \\rangle=0.53$) in comoving scales $\\rc\\approxlt 40-80\\hmpc$. No significant detection has been found in the DRS sample ($\\langle z\\rangle=1.4$) or in the RIXOS sample at small scales. Both tests are model--independent.  Quantitative measurements of the clustering and its evolution have been obtained from the integrated spatial correlation function (essentially the number of pairs with comoving separations smaller or equal than $\\rc$), assuming Poisson statistics, and a power--law shape for the correlation function, with slope $\\gamma=1.8$ and correlation length $r_0$.  Combining the limits from RIXOS ($\\rc\\approxlt 6\\hmpc$ and $\\rc\\approxlt 80\\hmpc$) with the DRS limits ($\\rc\\approxlt 6\\hmpc$), we obtain: $1.5 \\approxlt r_0\\approxlt 3.3\\hmpc$ for comoving clustering, $1.9 \\approxlt r_0\\approxlt 4.8$ for stable clustering and $2.2 \\approxlt r_0\\approxlt 5.5$ for linear evolution of clustering. Using $q_0=0$ we obtain $\\sim$20--30 per cent higher limits.  These results are compatible with the clustering properties of `normal' galaxies, but would imply weaker clustering in X--ray selected AGN than in optically--UV--selected ones. Our results do not support an increase in the clustering amplitude with redshift. However, we do not have many sources above $z>2.5$, so we cannot rule out neither support the strong clustering above that redshift, found by Stephens et al. (1997).   AGN (and/or NELGs) clustered like the sources studied here could produce most of the hard and soft X--ray background without exceeding the observed limits on its Auto Correlation Function.  The lack of very close pairs in RIXOS implies that the ratio of bias parameters between X--ray selected AGN and {\\sl IRAS} galaxies (where there is evidence that biasing is small, $b_I\\approx 1$) is $b_X/b_I \\approxlt 0.8-1.7$, somewhat smaller than previous results.     A 3 sigma detection at comoving separations $\\rc\\approxlt10\\hmpc$ would necessitate $\\sim 100-110$ `RIXOS--like' fields for comoving evolution and $r_0=6\\hmpc$. This is especially relevant for future X--ray survey identification programmmes (such as the XMM--XID) with a view to detect AGN clustering. It is much more efficient to use a compact connected area (circular or elongated) for this purpose, rather than serendipitous pointings distributed all over the sky.  In this sense, the \\ROSAT{ } All Sky Survey identification programmes will be very useful to provide a local `anchor' for X--ray AGN clustering studies (e.g. Zickgraf et al., 1998).    \\medskip \\noindent {\\bf ACKNOWLEDGEMENTS} \\medskip  \\noindent Special thanks are due to Omar Almaini and his colleagues for providing us with the positions, fluxes and redshifts of the DRS AGN. FJC and XB thank the DGES for financial support, under project PB95-0122. XB acknowledges sabbatical support at Cambridge to the DGES under grant PR95-490. ACF and RGM thank the Royal Society for support. This research has made use of data obtained from the Leicester Database and Archive Service at the Department of Physics and Astronomy, Leicester University, UK. We thank the Royal Society for a grant to purchase equipment essential to the RIXOS project.  {} {} {} {} {} {} {}"
    },
    "9804/astro-ph9804122_arXiv.txt": {
        "abstract": "We present the first fully relativistic investigation of matter fallback in a supernova. We investigate spherically symmetric supernova fallback using a relativistic radiation hydrodynamics Lagrangian code that handles radiation transport in all regimes. Our goal is to answer the fundamental question: did SN1987A produce a black hole and, if so, when will the hole become detectable ? We compute the light curve, assuming that a black hole has been formed during the explosion, and compare it with the observations. Our preliminary calculations lack radioactive energy input and adopt a very simple chemical composition (pure hydrogen). As a result, our computed models cannot fit the observed data of SN1987A in detail. Nevertheless, we can show that, during the first hours, the accretion flow is self--regulated and the accretion luminosity stays very close to the Eddington limit. The light curve is completely dominated, during the first few weeks,  by the emission of the stellar envelope thermal energy, and  resembles that obtained in ``standard'' supernova theory. Only long after hydrogen recombination takes place is there even a chance to actually detect radiation from the accreting black hole above the emission of the expanding envelope. The presence of a black hole is thus not inconsistent with observations to date. Because of the exponential decay of the $^{44}$Ti radioactive heating rate, the date of the emergence of the black hole is not very sensitive to the actual parameters of the models and turns out to be about 1000 years. The bulk of the emission then is expected to be in the visible band, but will be unobservable with present instrumentation. We discuss the implications of our results in connection with the possible emergence of a black hole in other supernovae. ",
        "introduction": "}  On the 23rd of February 1987, Shelton and Jones announced the discovery of a supernova in the Large Magellanic Cloud, SN1987A \\markcite{Kunkel et al. 1987}(Kunkel {\\it et al.\\/} 1987). Historically, this was the brightest supernova observed after that recorded by Kepler in 1604 (SN1604). It was also the first supernova to be observed in every band of the electromagnetic spectrum and the first detected through its initial burst of neutrinos. Its relatively close distance ($\\sim$ 50 Kpc) has offered a unique opportunity to observe a supernova in great detail and with a variety of detection techniques. [For extensive reviews on SN1987A see e.g. \\markcite{Arnett et al. 1989}Arnett {\\it et al.\\/} (1989), \\markcite{McCray 1993}McCray (1993) and references therein.] Since its appearance, SN1987A has confirmed many aspects of the theory of type II supernovae (see e.g. \\markcite{Dar 1997}Dar 1997). As predicted, most of the energy ($\\sim 10^{53}$ erg) was released in $\\sim$ 10 s by the cooling of the protoneutron star and was radiated in form of neutrinos, detected by the Mont Blanc, Kamiokande, IMB and Baksan underground detectors. The strong shock produced by the core bounce heated and pushed outwards the stellar envelope, which subsequently emitted its internal energy as photons: the integrated light emission was $\\sim 10^{49}$ erg, while the kinetic energy of the expanding shells was $\\sim 10^{51}$ erg. The light curve can be fitted quite well by the ``standard'' type II supernova theory (see e.g. \\markcite{Woosley 1988}Woosley 1988; \\markcite{Arnett 1996}Arnett 1996). After the emergence of the shock wave, the emission is dominated by the diffusion luminosity of the expanding stellar envelope during the first 30--40 days. At $\\sim 40$ days, the hydrogen envelope starts to recombine and a large amount of internal energy is released by the inward motion of the recombination front. After most of the envelope has recombined at day $\\sim 160$, the light curve starts to be dominated by the radioactive decay of heavy elements synthesized during the explosion. Some predictions of the ``standard'' theory were not confirmed and a number of key issues remain to be answered, such as the nature of the supernova progenitor, the explosion mechanism and, perhaps most important, the nature of the compact remnant left over after the explosion.  Stellar evolution calculations show that stars with a main sequence mass in the range 8--19 $M_\\odot$ finish their lives with a compact core of $\\sim 1.4 M_\\odot$. As shown by \\markcite{Woosley \\& Weaver 1995}Woosley \\& Weaver (1995), for these stars the amount of material that falls back toward the core in the aftermath of its collapse is negligible, so it is likely that they give birth to neutron stars. Stars with main sequence mass larger than $M_c \\sim$ 25--30 $M_\\odot$ have more massive cores ($\\simeq 2 M_\\odot$) and undergo accretion of a significant amount of matter, so that the remnant mass left over after the explosion is larger than $3 M_\\odot$, probably leading to the formation of a black hole. The value of $M_c$ is somewhat uncertain and depends quite sensitively on the mass loss prior to the explosion (reduced core mass) and on the explosion energy (amount of fall back). The fate of stars in the intermediate range 19--25 $M_\\odot$ is far less obvious. At the end of their evolution, they have core masses around 1.6--1.8 $M_\\odot$ and a variable amount of matter, 0.1--0.3 $M_\\odot$ may fall back because of the hydrodynamic interaction of the outgoing supernova shock with the expanding envelope (\\markcite{Woosley \\& Weaver 1995}Woosley \\& Weaver 1995). Typically, these stars leave a central compact remnant of 1.7--2.1 $M_\\odot$ (baryonic mass). Then their fate depends critically on the equation of state at nuclear matter densities that fixes the maximum mass $M_{crit}$ above which no stable neutron star configuration can exist. Adopting the parametrization of the nuclear matter equation of state by \\markcite{Lattimer \\& Swesty 1991}Lattimer \\& Swesty (1991), $M_{crit}$ turns out to be $M_{crit} \\lesssim 2.5 M_\\odot$ (gravitational mass). In this case, it would be unlikely that these stars could form a black hole. However, \\markcite{Thorsson, Prakash \\& Lattimer 1994}Thorsson, Prakash \\& Lattimer (1994) have proposed a much softer equation of state as a consequence of $K^{-}$ condensation. They find a maximum mass $M_{crit} \\simeq 1.5 M_\\odot$ (gravitational mass). Recently, modern models of nuclear forces and accurate Monte Carlo modeling of nucleon interactions indicate that $M_{crit} = 1.8$--2.2 $M_\\odot$ (\\markcite{Pandharipande 1997}Pandharipande 1997). Based on the findings of \\markcite{Thorsson, Prakash \\& Lattimer 1994}Thorsson, Prakash \\& Lattimer (1994), \\markcite{Brown \\& Bethe 1994}Brown \\& Bethe (1994) and \\markcite{Woosley \\& Timmes 1996}Woosley \\& Timmes (1996) have proposed different scenarios that may lead to the formation of a black hole after the explosion of a 19--25 $M_\\odot$ star. According to \\markcite{Brown \\& Bethe 1994}Brown \\& Bethe (1994), after core bounce a protoneutron star would form that could remain stable for about $\\sim 12$ s (just the time necessary to release its internal energy in form of neutrinos) and then would collapse to a black hole, as early suggested by \\markcite{Wilson et al. 1986}Wilson {\\it et el.\\/} (1986) and \\markcite{Woosley \\& Weaver 1986}Woosley \\& Weaver (1986). \\markcite{Woosley \\& Timmes 1996}Woosley \\& Timmes (1996) point out that the formation of a black hole could be driven by the matter which falls back in the first few hours after the explosion, as already emphasized by \\markcite{Colgate 1971}\\markcite{Colgate 1988}Colgate (1971, 1988) and \\markcite{Chevalier 1989}Chevalier (1989).  Interestingly, the mass of the progenitor of SN1987A ($M =$ 18--21 $M_\\odot$) falls in the range of masses where the outcome of core collapse is uncertain. So, from a theoretical point of view, we do not know if a neutron star or a black hole has been formed during the explosion. The main uncertainties are the value of the critical mass above which no stable neutron star can exist and the amount of matter that falls back during the explosion. If there was no fallback in SN1987A, the emission from a radio pulsar most probably would not be detectable directly for a very long time, unless there are holes in the outflow. Observations also rule out the presence of any optical pulsar. The energy deposition from a pulsar could alter the light curve, though; observations to date limit any energy input from a pulsar to be lower than the contribution from radioactive decay ($\\sim 10^{36}$ erg s$^{-1}$, implying $B/P^2 \\sim 10^{14}{\\rm G\\,s^{-2}}$ for magnetic dipole emission by a pulsar with field strength $B$ and rotation period $P$). However, high energy X-rays from a pulsar radiating a hard spectrum should have become observable in a few months (\\markcite{McCray, Shull and Sutherland 1987}McCray, Shull and Sutherland 1987) although the 1--3 keV thermal X--ray emission from the cooling of a neutron star would be attenuated for about 100 yr (\\markcite{Chernoff, Shapiro and Wasserman 1989}Chernoff, Shapiro and Wasserman 1989). Such hard X-rays were never observed.  If a compact object of any kind is present, it is probably accreting from the progenitor stellar material, since the inner part of the ejecta have velocity smaller than the escape velocity. During the first phases after the explosion, the accretion rate should be as high as $10^8$ in units of the Eddington accretion rate (\\markcite{Chevalier 1989}Chevalier 1989). On the other hand, as mentioned above, the bolometric light curve observed to date (ten years after the explosion) can be explained by the ``standard'' theory of Type II supernovae and does not show any feature related to the possible presence of a central accreting compact object. Using analytic and numerical methods, \\markcite{Chevalier 1989}Chevalier (1989) and \\markcite{Houck \\& Chevalier 1991}Houck \\& Chevalier (1991) have investigated the problem of fall back onto a neutron star and found that, for radiation pressure--limited accretion, the luminosity emitted after 3--4 years should approach the Eddington value, $L_{Edd} \\sim 10^{38}$ erg s$^{-1}$ for gas opacity dominated by electron (Thomson) scattering. Since at $\\sim 1500$ days the bolometric luminosity of SN1987A was $\\sim 10^{37}$ erg s$^{-1}$, they argue that either the accreting envelope becomes dynamically unstable or SN1987A contains a black hole. On the other hand, \\markcite{Chen \\& Colgate 1995}Chen \\& Colgate (1995) noted that, at the typical temperature and density of the ejecta after few years, the opacity in the accreting gas enriched of heavy elements is much larger than the Thomson opacity. So, they conclude that a neutron star accreting at the actual (non--Thomson) Eddington limit may be present. Furthermore, if the value of the critical mass falls in the range 1.8--2.2 $M_\\odot$, as seems to be indicated by recent developments in the nuclear many--body theory (\\markcite{Pandharipande 1997}Pandharipande 1997), it appears that the amount of fallback is a key element in determining whether or not a stellar black hole has formed. SN1987A may or may not harbor a black hole but observations alone have not resolved this issue yet. It is therefore important to determine what extra luminosity an accreting central component would produce and when its presence might be discernible. The amount of fallback depends on the main parameters of the explosion while the visibility of the remnant on the detailed hydrodynamical evolution of the flow and of the radiation field.  This is the first of a series of papers in which we address a number of relevant questions: can we infer from the light curve whether a stellar black hole has formed in the aftermath of a supernova explosion ? Did SN1987A produce a neutron star or a black hole ? We will consider these issues by assuming that either a black hole or a neutron star has formed and then will compute the resulting light curve varying the main parameters of the explosion. Our focus will be to isolate the contribution from the central compact object.  In this paper we present the first self--consistent, fully relativistic investigation of supernova fallback in presence of a black hole. It is the relativistic generalization of the work by \\markcite{CSW}Colpi, Shapiro \\& Wasserman (1996, hereafter CSW) in which the hydrodynamical properties of a fluid accreting onto a central remnant from an initially expanding cloud were first explored. Here, we include self-consistently the transfer of radiation. Due to the complexity of the problem we will not attempt to investigate a model with realistic chemical composition (that will be presented in a forthcoming paper) but, instead, we will consider a number of simpler preliminary scenarios to guide our understanding. We will show that, during the first hours, the accretion flow is self--regulated and the accretion luminosity stays very close to the Eddington limit. However, during the first few weeks the light curve is completely dominated by the emission of the stellar envelope internal energy, giving a light curve that resembles that obtained in the ``standard'' supernova theory. Although in this paper we apply our calculation specifically to SN1987A, our results have general validity and can be used to study the light curve and the observational signatures of black holes in other supernovae.  We have investigated spherically symmetric supernova fallback using a general--relativistic, radiation hydrodynamic Lagrangian code. Our code can handle the transfer of radiation from the first phases immediately after the supernova explosion (when photons diffuse through a high temperature, expanding cloud) to the late evolutionary stages (when the hydrogen envelope has recombined and most of the ejecta are transparent). The main challenge we have confronted is the enormous dynamic range in the relevant physical timescales entering the problem. To this end, we have implemented a multiple timestep procedure that allows us to integrate different radial regions at different rates (a primitive ``temporal adaptive mesh''). This technique enables us to speed up the calculation by almost an order of magnitude.  The plan of the paper is the following: in Section \\ref{sec2} we present the general relativistic equations of radiation hydrodynamics used in the calculation. The finite difference form of these equations is presented in the Appendix. Section \\ref{sec3} describes the numerical method used to solve the equations of radiation hydrodynamics and the boundary conditions adopted. In Section \\ref{sec4} we present in detail the multiple timestep procedure employed in our calculation. Section \\ref{sec5} illustrates the set up of the initial conditions. In Section \\ref{sec6} we present the numerical results for various simplified examples and test cases. Section \\ref{sec7} is devoted to the semi--analytic calculation of the late--time light curve. Finally Section \\ref{sec8} contains a discussion of our preliminary results in connection with SN1987A and with the diagnosis of black holes in supernova explosions. ",
        "conclusions": "}   The computed light curves of model IVa and IVb show that, prior to recombination, the release of internal energy from the expanding stellar envelope dominates over accretion. Although during the first few hours the diffusive accretion luminosity can be very close to the Eddington limit, the diffusive and advection luminosities of the hot outer layers are largely superEddington. As shown by model IVb, during recombination the computed light curve shows the typical behavior observed in Type II--Ib supernovae with low envelope mass. The luminosity at peak is comparable to that of SN1987A because the initial sound velocity of the hydrogen envelope is similar (see Tables \\ref{tab1} and \\ref{tab2}). The high expansion velocity and low envelope mass of model IVb reduce all of the relevant timescales (see Table \\ref{tab1}) so that recombination takes place at about 10 days after the onset of the explosion, instead of the 40--50 days inferred for SN1987A (see e.g. \\markcite{Woosley 1988}Woosley 1988). Nevertheless, our numerical results clearly indicate that no distinct signature in the light curve is found that could reveal the accreting black hole during this initial evolutionary phase. Only after the sharp decrease in luminosity due the clearing of the envelope by recombination, the accreting black hole starts to power the bolometric light curve of the computed models and $l$ decreases as $t^{-25/18}$ (see equations [\\ref{dustmdot}] and [\\ref{lumblon}]).  Slightly after maximum, SN1987A was powered by the radioactive decay of $^{56}$Co and $^{57}$Co and, at present, its light curve is consistent with emission from the decay of $^{44}$Ti (\\markcite{Suntzeff 1997}Suntzeff 1997; see  Figure \\ref{fig14}). Can we discern the luminosity emitted by the accreting  black hole above the contribution resulting from $^{44}$Ti ? As already noted, the accretion history is very sensitive to the value of the sound speed (see equation [\\ref{bondimass}]). Depending on the actual value of $c_{s,0}$ in the inner part of the ejecta surrounding the compact remnant (the mantle), the accretion rate and, in particular, the total accreted mass vary significantly. An upper limit to the accretion luminosity can be inferred taking the initial parameters of the mantle quoted by \\markcite{Chevalier 1989}Chevalier (1989; see Table \\ref{tab2}). In this case, $c_{s,0} \\simeq 3 \\times 10^7$ cm s$^{-1}$ and $t_{a,0} \\sim t_0$, giving substantial fallback. The mass accretion rate can be computed using equation (\\ref{dustmdot}) and the luminosity according to equation (\\ref{lumblon}). Thus, we have \\begin{equation} \\label{upper} l \\simeq 3\\times 10^{-7}(\\rho_0 \\, t_0 \\, c \\, k_{es})^{5/6}\\left ({t\\over t_0}\\right)^{-25/18}. \\end{equation} Following \\markcite{Chevalier 1989}Chevalier (1989), we adopt $\\rho_0\\, t_0^3\\approx 10^9$ cgs and $t_0\\approx 7000$ seconds, yielding a maximum accretion luminosity \\begin{equation} \\label{lumfall} l \\simeq {8\\times 10^{-3}\\over (M_{bh}/M_\\odot)^{1/3}[t({\\rm years})]^{25/18}} \\, . \\end{equation} Equation (\\ref{lumfall}) implies $L\\simeq 5\\times 10^{34}{\\rm erg\\,s^{-1}}$ after $t\\approx 10$ years, which is well below the present day bolometric luminosity of the remnant ($\\sim 10^{36}{\\rm erg\\,s^{-1}}$; Suntzeff 1997) and also smaller than the luminosity estimated to result from radioactive decay. Thus, there is no observation that rules out the possibility that a black hole resides inside the SN1987A remnant. At very late times, when $\\dot m \\la 0.03$, the corresponding upper bound would be \\begin{equation} \\label{late} l \\simeq 2\\times 10^{-8}(\\rho_0 \\, t_0 \\, c \\, k_{es})^2 \\left ({t\\over t_0}\\right )^{-10/3}, \\end{equation} or, using Table \\ref{tab2}, \\begin{equation} \\label{afterti} l = {7\\times 10^{-8}\\over [t({\\rm 10^3 years})]^{10/3}}. \\end{equation} In Figure \\ref{fig14} we plot the late--time light curve of SN1987A calculated using equations (\\ref{lumfall}) and (\\ref{afterti}). As radioactive decay plummets at around % 270--2700 years, it is clear that the black hole would only appear after about 900 years irrespective of the detailed numbers. After this time has elapsed, the luminosity of the remnant would be $\\sim 10^{32}{\\rm erg\\,s^{-1}}$, too dim to be detectable using present technology, and possibly even a challenge for our distant descendants, as it would be hard to distinguish from the multitude of low mass stars crowding its field in the Large Magellanic Cloud.  Since we have carried out a frequency--integrated calculation, we can only estimate the emission properties from the effective temperature and the photospheric gas temperature. During the optically thick accretion phase (when $\\dot {m}>0.03$), the numerical results show that $T_{eff}$ and $T_{ph}$ are very close and have a value  $\\sim$ 8000--10000 $K$. This corresponds to a mean photon energy in the few eV range, for which the bulk of the emission is expected in the visible band. Since $T_{eff}$ and $T_{ph}$ are close together the continuum spectrum should be roughly a blackbody. This is not surprising because the emission processes are thermal and the gas in the inner accreting region is optically thick. Only at very late times ($t \\gtrsim 2500$ yr), the accretion flow becomes optically thin and the gas temperature reaches very high values ($10^8$--$10^9 K$; this regime has not been treated in the present numerical computation). So, at the time of emergence, the inner part of the accretion flow will be still optically thick and will be emitting roughly a black body spectrum peaked in the visible band.   A number of interesting issues can be addressed within the framework presented in this paper. Can the presence of a stellar black hole in SN1987A be deduced from the excess accretion luminosity it produces ? What is the observational signature of the explosion of more massive progenitors (those with $M\\sim 30 M_{\\odot}$) on the light curve of the remnant ? Woosley, \\& Timmes (1996) suggested that a distinguishing signature would be a bright plateau that steeply falls to very low or zero luminosity, similarly to what found in our models.  A number of  assumptions limit the applicability of our results, amongst, the hypothesis of spherical symmetry and the simple chemical composition considered. As a result of an asymmetry in the explosion mechanism, black holes may come to birth with significant intrinsic velocities (a few $\\times \\, 100$ km $^{-1}$). This is certainly observed for pulsars whose mean birth speed is around 250--300 km s$^{-1}$ \\markcite{Hansen \\& Phinney 1997}(Hansen \\& Phinney 1997; \\markcite{Cordes \\& Chernoff 1997}Cordes \\& Chernoff 1997). However, since the initial sound speed of the ejecta can be of the order of $10^3$ km s$^{-1}$, the accretion will proceed essentially subsonically. The central compact remnant tends to comove with gas expanding at its velocity and, if the ambient density is uniform, the conditions are independent of position within the core, as in the cosmological Hubble flow. Thus, even a relatively large intrinsic velocity should not strongly modify the results presented here. Deviations from spherical symmetry can arise if infalling matter has even a small amount of specific angular momentum, which presents a barrier to spherical infall (\\markcite{Chevalier 1996}Chevalier 1996). In the case of SN1987A, the fact that the observed neutrino burst is consistent with implosion models without rotation \\markcite{Burrows 1988}(Burrows 1988) seems to indicate that angular momentum of matter initially close to the black hole was not very important (see also \\markcite{Chevalier 1996}Chevalier 1996). In this case, the accretion should have proceeded almost radially, although mixing of outer mantle material with significant angular momentum into the inner region could modify our picture.  A more realistic model with shell--like chemical composition and heating due to radioactive decay will enable us  to address more quantitatively the problem of the visibility of black holes in otherwise successful supernova. This analysis will also set the basis for studying fallback onto a nascent neutron star in order to verify whether major outflows establish that reverse the process, during the decline of the supernova light curve."
    },
    "9804/astro-ph9804136_arXiv.txt": {
        "abstract": "In a recent paper, Hawkins (1997) argues on the basis of statistical studies of double-image gravitational lenses and lens candidates that a large population of dark lenses exists and that these outnumber galaxies with more normal mass-to-light ratios by a factor of 3:1. If correct, this is a very important result for many areas of astronomy including galaxy formation and cosmology. In this paper we discuss our new radio-selected gravitational lens sample, JVAS/CLASS, in order to test and constrain this proposition. We have obtained ground-based and HST images of all multiple-image lens systems in our sample and in 12 cases out of 12 we find the lensing galaxies in the optical and/or near infrared.  Our success in finding lensing galaxies creates problems for the dark lens hypothesis. If it is to survive, ad hoc modifications seem to be necessary:  only very massive galaxies ($M\\ga9\\times10^{11}{\\rm M}_{\\odot}$) can be dark, and the cutoff in mass must be sharp. Our finding of lens galaxies in all the JVAS/CLASS systems is complementary evidence which supports the conclusion of Kochanek et al. (1997) that many of the wide-separation optically-selected pairs are physically distinct quasars rather than gravitational lens systems. ",
        "introduction": "Gravitational lensing is an important phenomenon because it probes mass distributions of galaxies independently of their optical luminosity. In particular, it provides a unique way to search for galaxies with very high mass-to-light ratios which are otherwise difficult to detect.  In a recent paper, Hawkins (1997) performs such a search, based on published data on eight two-image gravitational lens systems and candidate lens systems with image separations $>2^{\\prime\\prime}$. Statistical arguments are presented that these are indeed genuine lens systems rather than chance associations of unrelated quasars. For example, several of the quasar pairs have more similar colours than one would expect to find in the general population of quasars, and in at least one case that the spectra are so similar that when one spectrum is divided by the other, the result is a constant ratio to within the noise (Hawkins et al. 1997). No lensing galaxies are found in six of the systems, resulting in mass-to-light ratios of up to 22000 in the most extreme case. The inference is that the lensing is done by underluminous `dark galaxies' with very substantial components of dark matter, with serious implications for cosmology as well as lensing studies. A further paper by Jimenez et al. (1997) discusses how such galaxies could be formed. However, Kochanek, Falco \\& Mu\\~{n}oz (1997) have argued against a lensing interpretation and for the hypothesis that the Hawkins lenses are binary quasar pairs, based on existing statistics of large separation lenses, because a population of wide-separation, optically-selected lenses should imply a significant, unseen, population of corresponding radio-selected lenses. They also discuss formation scenarios for such binary quasar pairs. In this paper we describe a new, well-defined, sample of gravitational lens systems selected from the JVAS/CLASS radio surveys.  We discuss the search for the lensing galaxies in the 12 systems found so far and assess the implications for the proposal that dark lensing galaxies exist. In Section 2 we give a brief description of JVAS/CLASS and list the confirmed lens systems. In Section 3 we summarize the observations, both with ground-based telescopes and with the HST, in which we detect lensing galaxies in all 12 systems. Finally, in Section 4, we present our conclusions and discuss further the question of whether the optically-selected pairs (Hawkins 1997) are lens systems, physically unrelated quasar pairs or related quasar pairs. ",
        "conclusions": "In our radio-selected lenses, even without the detection of the lensing galaxy, there can be no doubt that one is dealing with gravitational lens systems, because of radio spectral index and radio/optical similarities of the images.  The results we have obtained suggest strongly that lensing galaxies have $M/L_{I}\\la50$ and thus are not dark.  Twelve out of twelve of the JVAS/CLASS lensed systems have $M/L_{I}$ ratios suggestive of normal luminous galaxies. Even if {\\em all} the remaining 5 candidates were lensed systems, and {\\em all} of these had no detectable lensing galaxy, this would still leave over 60\\% of lens systems with detected, and relatively normal, lensing galaxies compared to $<$25\\% reported by Hawkins (1997).  There is a possible explanation for the difference between our results and those of Hawkins. This is that the Hawkins sample has a fundamentally different distribution of image separations, as it contains only systems with separations greater than 2$^{\\prime\\prime}$, which, since the separation $s$ is a function of lens mass $M$ (e.g. Schneider et al. 1992), corresponds to relatively massive lensing galaxies  ($\\ga9\\times10^{11}{\\rm M}_{\\odot}$, see Table 1). Could only massive galaxies be dominated by dark matter?   Let us concentrate, therefore, on radio lensed systems with separation $>$2~arcsec. We will further restrict our consideration to those incontrovertible lens systems in which optical counterparts to the radio components have been detected.  There are two such systems with separation $>$2~arcsec in the JVAS/CLASS sample\\footnote{We do not include B2114+024 as there is no detection of optical counterparts to any of the radio images.}. There are two further known radio-selected systems with $s>2^{\\prime\\prime}$; B0957+561 (from Hawkins' sample) and B2016+12 (Lawrence et al. 1984). In all four of these normal (not dark) lensing galaxies are detected. Thus amongst {\\em confirmed} lens systems with separations $>2^{\\prime\\prime}$ there is no evidence for dark lenses.   There is no reason why any bias should be introduced by the use of radio selected systems; it is simply a way of securely identifying lens systems without the necessity of identifying a lensing galaxy. The discrepancy between the results for radio-selected lens systems, even for large separation lenses, and the mostly optically-selected sample of Hawkins (1997) is therefore puzzling.  We note, however, that Hawkins only considers systems with 2 images. Among the radio-selected systems there are many 4-image systems. Could the dark lenses be found only in the 2-image systems?  Kochanek et al. (1997) pointed out that in general, one would expect a greater preponderance of 4-image systems where the lensing mass distribution, particularly the halo, is more elliptical (e.g. King \\& Browne 1996), and that one might therefore postulate that `dark lenses' have spherically symmetric mass distributions.  This, however, fails to account (Kochanek et al. 1997) for why the existing radio-selected two-image systems do not have dark lenses, as well as for why we find no two-image systems with dark lenses in JVAS/CLASS.   In summary, there are two conditions that must be met in order to reconcile our observations with the idea of dark lenses. First, in order for us to find no dark lenses in the JVAS/CLASS sample, there must be a sharp cutoff such that all dark lenses must have masses of $\\ga9\\times10^{11}{\\rm M}_{\\odot}$. Second, in order to explain the detection of lenses in the radio lens systems with separations $>$2~arcsec (mostly quads), dark lenses should have much rounder mass distributions than the luminous lenses (Kochanek et al. 1997). The alternative to these two conditions is to suppose that most, or even all, of the systems discussed by Hawkins in which no lensing galaxy is found, are indeed not lens systems. This is despite the fact that in each individual case we find the arguments persuasive that the quasar pairs are not random associations.  However, if there are no dark lenses, the frequency of close physically associated pairs of quasars has hitherto been greatly underestimated."
    },
    "9804/astro-ph9804300_arXiv.txt": {
        "abstract": "We compute the luminosity function of galaxies in groups. The data consists in two samples of groups of galaxies selected in distance and redshift space comprising a total number of 66 groups. The assignment of galaxies to the groups were derived from the Edinburgh-Durham Southern Galaxy Catalog, considering a limiting apparent magnitude $m_{lim}=19.4$.  We find a flat faint end of the mean galaxy luminosity function in groups in contrast to the mean galaxy LF in clusters where a large relative number of faint galaxies is present. We find that a Schechter function with parameters $M^{*}\\simeq -19.6\\pm 0.2,$ $\\alpha \\simeq -1.0\\pm $ $0.2$ provides a satisfactory fit to the galaxy LF of our total sample of groups in the range of absolute magnitudes $-22<M<-16$. ",
        "introduction": "Several works have attempted the identification of groups of galaxies in 2 dimensions using catalogs with angular positions (see for instance \\cite{tur76}). A major problem with such identifications is the low density contrast of these systems and therefore redshift surveys are essential to provide reliable physical groups. Nevertheless, most of redshift surveys are limited to bright apparent magnitudes allowing  only the identification of the brightest members. This fact make difficult the  determinations of the faint end of the luminosity function (hereafter LF) of galaxies in groups. The galaxy LF in clusters and nearby groups has been studied by several authors (\\cite{oem74}, \\cite{dre78}, \\cite{oeg86}, \\cite{oeg87}, \\cite{gud91}, \\cite{wil90}, \\cite{gar91}, \\cite{fer91}, \\cite{gar92}, \\cite{lop96}, \\cite{gai97}, \\cite{val97}, \\cite{tren97}). Of particular importance is the possible universal character of the galaxy LF early suggested by \\cite{abe62} and \\cite{abe75} since it may serve to pose constraints to models of galaxy formation as well as testing the importance of environment in the related astrophysical processes. Interactions play a fundamental role in galaxy evolution (see for instance \\cite{pos84}) being particularly relevant in small galaxy associations. Moreover, in contrast to clusters, groups have a low galaxy velocity dispersion, gas density and temperature which may cause a significantly different evolution of galaxies in groups and clusters. Groups of galaxies are therefore suitable systems to provide a useful observational insight on different physical phenomena related to environment.  The field galaxy LF has been determined in several works, see for instance \\cite{lov92}, and more recently \\cite{mar94} and \\cite{lin96}. These authors provide Schechter function fits (\\cite{sch76}) and in spite of discrepancies in the value of the parameter $M^{*}$, both determinations are consistent with a flat LF at the faint end, $\\alpha \\simeq -1.$ On the other hand, several studies of the galaxy luminosity function in clusters show a much steeper galaxy LF at the faint end, $\\alpha \\leq -1.4$ (e.g. \\cite{fer91}, \\cite{val97}, \\cite{tren97}).  The universality of the galaxy LF in clusters has been seriously questioned by \\cite{lc95} and  \\cite{lc97}. From their sample of 45 Abell clusters with z$<$0.14 39 clusters show an increase of the fraction of dwarf galaxies consistent with a double Schechter fit with $\\alpha \\simeq 1$ and $-2<\\alpha <-1.4$ and only 7 can be suitable fitted with a single Schechter fit with $\\alpha \\simeq 1$. This set of clusters with flat LF are characterized by an evolved morphology of the cD type and are on average massive and gas rich. Nevertheless, \\cite{driver} and \\cite{tren97} find rising LF in rich clusters which suggest the complexity of this phenomenon. In spite of the reported differences of the galaxy LF in clusters the results show agreement with the mean galaxy LF in clusters derived by \\cite{val97} which rises at faint magnitudes consistent with a single Schechter fit with $\\alpha \\simeq -1.4$. The tendency of poorer clusters to show a flatter galaxy LF ($\\alpha \\simeq -1.2$) than richer clusters ($\\alpha \\simeq -1.5$)  was obtained by \\cite{val97} dividing their sample in two equal number of clusters. These results and \\cite{lc97} findings would not be necessarily inconsistent given that the flat LF clusters ($\\alpha \\simeq -1.0$) observed correspond to a small fraction ($\\simeq 20\\% $) of dynamically relaxed systems where evolutionary process have an important role. Other works provide useful studies of the galaxy LF in different environments. For instance, \\cite{gai97} develop an important work on the luminosity function of galaxies in clusters. Nevertheless, the background subtraction in this work is taken at 0.5 Mpc of the cluster center which may contribute significantly to the resulting flat LF of the analysis given the expected large contamination of the background by cluster members.  Given the observed differences between cluster and field mean galaxy luminosity functions we analyze in this work the galaxy LF in moderate density associations. In order to provide a suitable study of the effects of environment we use the same statistical procedures as in \\cite{val97} avoiding possible systematic effects due to the use of different techniques. In section 2 we describe the group and galaxy data used. Section 3 gives a description of the methods of analysis adopted, the main results obtained and error estimates through a Monte-Carlo method. In Section 4 we present the main conclusions. ",
        "conclusions": "We have applied statistical analyses to calculate the mean galaxy luminosity function in groups. We have used COSMOS survey galaxies and group catalogs with either available distances or redshifts. Schechter functions provide good fits to the galaxy LF of the sample of groups analyzed. The best fitting parameters for our total sample of 66 groups are $\\alpha =-1.0\\pm 0.2$, $M^{*}=-19.6\\pm 0.2$. This value of the $\\alpha $ parameter contrasts with that corresponding to a single Schechter fit in clusters $\\alpha \\approx -1.4$ (see for instance \\cite{fer91}, \\cite{val97}) indicating an important relative excess of faint galaxies in clusters. The dependence of the galaxy LF on cluster richness was explored by \\cite{val97} finding that poor clusters show a flatter galaxy LF than more massive systems ($\\alpha \\simeq -1.2$ in contrast to $\\alpha \\simeq -1.5$). \\cite{lc97} results show a flat galaxy LF ($\\alpha \\simeq -1.0$) in a subsample of dynamically relaxed systems  which may  be a strong indication that evolutionary process may disrupt dwarf galaxies contributing to the formation of cD haloes in these systems. However, \\cite{driver} in their analyses of the galaxy LF of the rich (R=3) cD-dominated cluster A963 a large excess of faint galaxies is found which indicates the non universality of this phenomenon. Similar large excess of faint galaxies are reported by \\cite{tren97} for 4 rich clusters (R=2-4) indicating that the dependence of the faint excess on different physical parameters of clusters deserve further analysis. The importance of the cluster merging history and the gas-dynamical evolution within these systems should be addressed before firm conclusions can be derived from the analysis of the observations.  Our comparison with \\cite{val97} results have the advantage that both have used similar methods of analysis, same galaxy catalog and background subtraction procedure and therefore it provides a reliable confrontation of the mean galaxy LF in clusters and groups. In spite of the fact that our sample galaxies correspond to moderate galaxy overdensities we find no evidence for an excess of faint galaxies relative to the field as it is found for the mean galaxy LF in clusters. Although galaxy mutual interactions are expected to play an important role in groups our results indicate that at these moderate galaxy density enhancements the galaxy luminosity function is not significantly different from the field."
    },
    "9804/astro-ph9804070_arXiv.txt": {
        "abstract": "We have discovered two simultaneous kHz quasi-periodic oscillations (QPOs) in the Z source GX\\,340+0 with the Rossi X-ray Timing Explorer. The X-ray hardness-intensity and color-color diagram each show a full Z-track, with an extra limb branching off the flaring branch of the Z. Both peaks moved to higher frequencies when the mass accretion rate increased. The two peaks moved from $247 \\pm 6$ and $567 \\pm 39$ Hz at the left end of the horizontal branch to $625 \\pm 18$ and $820 \\pm 19$ Hz at its right end. The higher frequency peak's rms amplitude (5--60 keV) and FWHM decreased from $\\sim 5\\%$ and $383 \\pm 135$ Hz to $\\sim 2\\%$, and $145 \\pm 62$ Hz, respectively. The rms amplitude and FWHM of the lower peak were consistent with being constant near 2.5 \\% and 100 Hz. The kHz QPO separation was consistent with being constant at $325 \\pm 10$ Hz. \\newline \\indent Simultaneous with the kHz QPOs we detected the horizontal branch oscillations (HBO) and its second harmonic, at frequencies between 20 and 50 Hz, and 38 and 69 Hz, respectively. The normal branch oscillations were only detected on the upper and middle normal branch, and became undetectable on the lower normal branch. The HBO frequencies do not fall within the range predicted for Lense-Thirring (LT) precession, unless either the ratio of the neutron star moment of inertia to neutron star mass is at least 4, $10^{45}g\\,cm^2/M_{\\odot}$, the frequencies of the HBO are in fact the sub-harmonic oscillations, or the observed kHz peak difference is half the spin frequency and not the spin frequency. \\newline \\indent During a 1.2 day gap between two observations, the Z-track in the hardness-intensity diagram moved to higher count rates by about 3.5\\%. Comparing data before and after this shift, we find that the HBO properties are determined by position on the Z-track and not directly by count rate or X-ray colors. ",
        "introduction": "GX\\,340+0 is a bright low-mass X-ray binary and a Z source (\\cite{hasinger}). The Z-track traced out in the X-ray color-color (CD) and hardness-intensity diagram (HID) is usually built up of three components, the horizontal branch (HB), the normal branch (NB), and the flaring branch (FB). The mass accretion rate ($\\dot{M}$) is thought to increase from a few tens of the Eddington mass accretion rate when the source is on the HB, to near-Eddington on the NB, and super-Eddington on the FB. On the HB and the upper part of the NB of GX\\,340+0, quasi-periodic oscillations (QPOs) with frequencies of 32--50 Hz, the horizontal branch oscillations (HBOs), were detected by Penninx et al. (1991). On the central part of the NB, QPOs with a typical frequency of 5.6 Hz, the normal branch oscillations (NBOs), were detected by van Paradijs et al. (1988). \\newline \\indent In four Z sources, Sco\\,X-1 (\\cite{vanderKlis1996a}; \\cite{michiel1997b}), GX\\,5-1 (\\cite{michiel1996b}), GX\\,17+2 (\\cite{wijnands1997b}), and Cyg\\,X-2 (\\cite{wijnands1997c}) two simultaneous kilo Hertz (kHz) QPOs have been observed, with frequencies reflecting flux changes on timescales close to the dynamical timescales near a neutron star. The higher frequency peak (henceforth called upper peak) has been proposed to reflect the Keplerian frequency of blobs of material orbiting around the neutron star, at a preferred radius in the disk, such as the sonic radius (Miller, Lamb \\& Psaltis 1998). The lower frequency peak (henceforth called lower peak) is the proposed beat frequency between the frequency of the upper peak and the spin frequency of the neutron star. The kHz QPO frequency is observed to increase with increasing $\\dot{M}$. Similar kHz QPOs are also found in atoll sources. The frequencies of the kHz QPOs are between 300 and 1200 Hz. (See van der Klis (1997) for a recent overview). \\newline \\indent In this Letter we report the discovery of two simultaneous kHz QPOs in the Z source GX\\,340+0. A preliminary announcement of these results was already made by Wijnands \\& van der Klis (1998b). ",
        "conclusions": "We have discovered kHz QPOs in the Z source GX\\,340+0. The frequencies of both the lower and upper peak increased as the source moved along the HB to the HB/NB vertex. The peak separation was consistent with being constant. This is similar to what was found in GX\\,17+2 (\\cite{wijnands1997b}), Cyg\\,X-2 (\\cite{wijnands1997c}), and GX\\,5-1 (\\cite{michiel1996b}; \\cite{wijnands1998c}). Only in Sco\\,X-1 (\\cite{michiel1997b}) and perhaps in the atoll source 4U 1608-52 (\\cite{mendez1998}) the peak separation was found to decrease with increasing $\\dot{M}$. If this decrease in Sco X-1 is in some way related to the source approching the Eddington critical luminosity (\\cite{white1997}; \\cite{milleretal}), this could mean that if we were able to detect the kHz QPOs in GX\\,340+0 and the other Z-sources (further) up the NB we might also see the peak separation decrease. \\newline \\indent The lower peak reaches the so far lowest ``kHz'' QPO frequency found in any low mass X-ray binary, 247 Hz. The maximum frequency reached by the upper peak, 820 Hz, is also rather low. Using the sonic-point model (\\cite{milleretal}) to explain the upper peak kHz QPO frequency, we found lower limits on the mass of the neutron star of $M_{neutr} > 0.6M_{\\odot}$ for a neutron star radius of 4M, and  $M_{neutr} > 1.2M_{\\odot}$ for a neutron star radius of 7M. (The implicit assumptions made in calculating these lower limits are: Schwarzschild spacetime metric, non-rotating radiating layer, and photons remove angular momentum from the disk material only in one scattering event.) We note that the frequencies of the kHz QPOs seem to confirm the idea (e.g. \\cite{kuulkersetal}) that GX\\,340+0 is more similar in appearance to GX\\,5-1 (\\cite{wijnands1998c}), than to Sco\\,X-1 (\\cite{michiel1997b}), and GX\\,17+2 (\\cite{wijnands1997b}) (frequencies higher than 1000 Hz). \\newline \\indent Any straightforward beat-frequency model predicts the upper peak to be narrower than, or as narrow as the lower peak. We find in GX\\,340+0 that the upper peak is usually broader than the lower peak. In fact, the upper peaks found in GX\\,340+0 were very broad with $\\frac{\\delta \\nu}{\\nu}$ up to 0.67 at $S_z=0.43$. Therefore a straightforward beat-frequency interpretation of the two peaks in GX\\,340+0 is difficult. In GX\\,340+0, GX\\,5-1 (\\cite{wijnands1998c}), and GX\\,17+2 (\\cite{wijnands1997b}) the upper peak FWHM decreases as a function of $S_z$. To explain this behaviour one could involve scattering in a rapidly variable medium (\\cite{mendez1998}). The effect of this scattering should become less as $\\dot{M}$ increases. \\newline \\indent Although the behaviour of the kHz QPOs as a function of inferred $\\dot{M}$ in GX\\,340+0 is similar to that found in other Z sources, the HBO behaviour is rather different. The FWHM of the HBO increased as the rms amplitude decreased, which is unlike what is seen in other Z sources. In Cyg\\,X-2 Wijnands et al. (1997c) found the FWHM to be well-correlated to the rms. In GX\\,17+2 the FWHM remained about constant while the rms amplitude decreased (\\cite{homan}). In previous observations of GX\\,340+0 (\\cite{penninx}, Table 2) the behaviour is consistent with our findings. In observations of Kuulkers \\& van der Klis (1996) the FWHM does not show a clear relation as a funtion of the rms amplitude. Exactly the same behaviour (as we reported for GX\\,340+0) was found in GX\\,5-1 (\\cite{lewin1992}; \\cite{wijnands1998c}). This suggests that an extra parameter in addition to $\\dot{M}$ determines the timing properties. Another indication of this is that there are differences between the Z sources in the position on the Z-track where the NBO are found. In GX\\,17+2 (\\cite{penninx1990}), and Sco\\,X-1 (\\cite{vanderKlis1987}) the NBO were detected up the FB, e.g. $S_z > 2$. In Cyg\\,X-2 they were detected up to the NB/FB vertex (\\cite{wijn1997a}), while in GX\\,340+0 and GX\\,5-1 the NBO were detected up to $S_z=1.8$. \\newline \\indent Recently it has been proposed that the HBO frequency may reflect the general relativistic Lense-Thirring (LT) precession frequency at the inner edge of a warped accretion disk (\\cite{stella}). If the upper peak is interpreted as a Keplerian frequency in the inner part of the disk, there is a quadratic relation between the LT precession frequency and the upper peak frequency. Applying this theory to our observations and assuming the peak difference to be the neutron star spin frequency, the LT predictions for the HBO frequency are too low by at least a factor of 2. The same discrepancy by a factor of 2 has been reported for the Z source GX\\,17+2 (\\cite{stella}). There are several ways out, but each appears to create serious difficulties. $I/M$ (with I the moment of inertia of the neutron star and M its mass) may be a factor of 2 higher than presently thought, however, no sensible equation of state appears to be able to produce such values of $I/M$. The kHz QPO peak difference may not be equal to the neutron star spin frequency but to half the spin frequency, however, at least in atoll sources this seems to contradict the results obtained from burst QPO (\\cite{strohmayer}). The HBO may not be the fundamental frequency but the harmonic of the predicted LT frequency, however, in none of the Z sources there is any evidence of excess power at frequencies half the HBO frequencies."
    },
    "9804/astro-ph9804246_arXiv.txt": {
        "abstract": "The dependence of the thermal pressure of hot galactic halos on a model parameter describing the frequency of major reheating episodes during galactic history is investigated. Pressure on the interface between interstellar medium and the halo gas is especially interesting, since empirical evidence here offers one of the simplest constraints on halo models. It is shown that two-phase model of Mo \\& Miralda-Escud\\'e is sufficiently robust with respect to uncertainties in the average interval between reheating. ",
        "introduction": "The existence of very extended galactic halos\\footnote{The term ``halo'' will be preferred over such frequently used terms as ``corona'' or ``envelope'', although it is implied that they all describe the same physical objects.} has been suspected for a long time (Spitzer 1956) on several grounds. In last two and a half decades, from the observations of {\\it Copernicus} and {\\it IUE\\/} satellites we have learnt that highly ionized gas exists in the halo of Milky Way to much larger scale-heights than it was previously assumed (for a review, see Savage 1988). So-called high-velocity clouds were discovered high above the plane of the Galaxy, and in other nearby spiral galaxies. X-ray observations showed the presence of vast quantities of hot gas in rich clusters of galaxies (Sarazin 1986), as well as in smaller compact groups (Saracco \\& Ciliegi 1995; Mulchaey et al. 1996). Cooling flows phenomena were observed in cluster members (Fabian 1994), as well as in individual elliptical galaxies (Nulsen, Stewart, \\& Fabian 1984). The extraplanar optical recombination emission was discovered at large galactocentric distances in several nearby galaxies (Donahue, Aldering, \\& Stocke 1995; Pildis, Bregman, \\& Schombert 1994). The depletion of globular cluster gas was interpreted as the consequence of ram-pressure stripping, and consequences were drawn thereof (Frank \\& Gisler 1976; Ninkovi\\'c 1985).  Especially important motivation comes from the QSO absorption line studies. It is now clear that, at least at low redshift, a bulk of both metal and Ly$\\alpha$ absorption lines arise in gas which is associated with luminous galaxies (Sargent, Steidel, \\& Boksenberg 1988; Bechtold \\& Ellingson 1992; Lanzetta et al. 1995; Bowen, Blades, \\& Pettini 1996; Chen et al. 1998). The idea that the QSO absorbers might arise in halos of normal galaxies originated with Bahcall \\& Spitzer (1969), but the quantitative models were rare (e.g. Bregman 1981) until recently, when vast accumulated data on absorber statistics enabled establishing strong empirical constraints on model parameters (see discussion in Mo \\& Morris 1994). Very important simple two-phase model of the galactic gaseous halos was put forward recently by Mo (1994), and Mo \\& Miralda-Escud\\'e (1996; hereafter MM96). Their work shows quite persuasively that formation of the halo structure as a natural consequence of the galaxy formation and subsequent evolutionary processess.  Formation of galaxies implies the collapse of the baryonic content of intergalactic space and its motion through the non-dissipative dark matter halos creating deep gravitational potential well. A halo of hot gas at virial temperature ($T_v \\equiv \\mu V_{\\rm cir}^2/2k$, where $\\mu$ is the average mass per particle, $k$\\/ is the Boltzmann constant and $V_{\\rm cir}$ is the circular velocity, which is assumed to parametrize the total mass of the galaxy) will form as the kinetic energy of the infalling material is thermalized in shocks of this primordial accretion. As the gas collapses, it is shocked and it subsequently cools (for an early treatment, see White \\& Rees 1978). Afterwards, similar situation (although presumably on smaller scale) arises in course of major galaxy mergers: here we also have a large mass of gas accreted into the system on a short timescale and substantial reheating. In the meantime, accretion of ambiental IGM is small or entirely negligible. When hot gas cools, it forms the {\\it cold phase}: photoionized clouds at $T_{\\rm cold} \\sim 10^4$ K, in pressure equilibrium with the hot halo, slowly sinking in the galactic gravitational potential. While cold phase may dominate the total gas mass, its filling factor is always small, and we shall not deal with it here separately.  The galactocentric radius $r_c$ where the cooling time of the hot gas is the {\\it cooling radius}. Quantitatively, it is necessary to satisfy \\begin{equation} \\label{minone} \\frac{5\\mu k T_v}{2 \\Lambda(T_v) \\rho(r_c)} = t_M, \\end{equation} where $\\Lambda(T_v)$ is the cooling rate at the virial temperature, and $\\rho(r_c)$ is the physical density of hot gas at the cooling radius. Inside it, the conditions of adiabaticity and hydrostatic equilibrium, lead to the equation of state $P \\propto \\rho^{5/3}$, where $\\rho$ is the density of the gas. The ratio of the total gas mass and the virial mass is set by the dimensionless parameter $f_g$. Isothermal potential of the dark matter creates the density profile \\begin{equation} \\label{zero} \\rho \\propto \\left( 1- K \\ln \\frac{r}{r_c} \\right)^{\\frac{3}{2}}, \\end{equation} where $K$ is the constant, in the simplest model equal to 0.8. Simple calculation shows that pressure at some distance $r$ will behave as $P/k \\propto (1-K \\ln \\frac{r}{r_c} )^{5/2}$.  We see that one of the model parameters is $t_M$, the average interval between reheating of the gas, the inverse of what we may call the reheating frequency. It is one of the most uncertain parameters. MM96 use \\begin{equation} \\label{one} t_M =\\frac{t}{1+\\Omega_0} \\end{equation} approximation, where $t$ is the age of the universe, and $\\Omega_0$ is the total cosmological density parameter. It is used quite successfully by MM96, although admitting that a detailed treatment is difficult (due not only to the uncertainty in $\\Omega_0$, but the power spectrum and physics of structure formation as well). While not questioning the goodness of this approximation, we would like to show that actual physical quantities resulting from their model are not much affected by the choice of $t_M$.  One of, from the physical point of view, most important predictions is the value of the thermal pressure of the hot halo at small galactocentric distances, where interface with the ISM occurs. Therefore, in this paper, we are investigating the changes in the pressure at this inner boundary of the hot halo, when the average time-scale for reheating of the gas is varied. This is a part of the more comprehensive study of the boundary conditions of the two-phase model, which is currently in progress (\\'Cirkovi\\'c 1998). ",
        "conclusions": "Major weaknesses of the theoretical framework this simple, manifest themselves clearly in the fact that mergers are considered only in terms of reheating through arising shocks, neglecting the mass input which undoubtedly occurs. On the other hand, without a detailed understanding of the underlying physics of merging events, it is difficult to proceed in that direction. It is clear, though, that if the bulk of the mass deposition occurs at small enough distances (compared to the virial radius before the merger), hot gas will be just a transient phase before most of the mass goes into cold, photoionized clouds. This can be estimated as follows: if $t_M$ is large compared to the timescale given by $\\tau=5 \\mu k T_0/2 \\Lambda(T_0) \\rho(r_0)$ (where $T_0=T_v(1-K \\ln \\frac{r_0}{r_c})$ and $\\rho(r_0)$ is given by the equation (\\ref{zero}), with $r_0$ being the characteristic radius for mass deposition), a quasi-stationary state is reached, and the discussion we have followed applies.  We have so far discussed only rehetaing in major mergers, which undoubtedly occurs from empirical evidence (Keel \\& Wu 1995, and references therein). Other conceivable reheating mechanisms (starbursts or switching on of a nuclear source) may be important in the inner part of the halo and should be considered in the course of a future work.  We have shown that simple model based on assumptions of adiabaticity and quasi hydrostatic equilibrium gives physically acceptable results for the pressure at the ISM-halo interface, results which are quite insensitive to the assumptions about reheating frequency $1/t_M$. Until our theoretical knowledge on merger frequencies (or other reheating mechanisms) improves, we are justified in using simple approximations like (\\ref{one}).  \\vspace{0.5cm}  The author is happy to acknowledge help of Dr. Hou Jun Mo, who kindly provided a cooling code, as well as useful discussion, and Mr. Branislav Nikoli\\'c whose hospitality and support were essential for completion of this work. Dr Luka \\v C. Popovi\\'c was, as usual, extremely helpful with a friendly advice and technical assistance."
    },
    "9804/astro-ph9804185_arXiv.txt": {
        "abstract": "We describe planetesimal accretion calculations in the Kuiper Belt. Our evolution code simulates planetesimal growth in a single annulus and includes velocity evolution but not fragmentation.  Test results match analytic solutions and duplicate previous simulations at 1 AU.  In the Kuiper Belt, simulations without velocity evolution produce a single runaway body with a radius $r_i \\gtrsim$ 1000 km on a time scale $\\tau_r \\propto M_0^{-1} e_0^{x}$, where $M_0$ is the initial mass in the annulus, $e_0$ is the initial eccentricity of the planetesimals, and $x \\approx$ 1--2. Runaway growth occurs in 100 Myr for $M_0 \\approx$ 10 \\mearth~and $e_0 \\approx 10^{-3}$ in a 6 AU annulus centered at 35 AU. This mass is close to the amount of dusty material expected in a minimum mass solar nebula extrapolated into the Kuiper Belt.  Simulations with velocity evolution produce runaway growth on a wide range of time scales.  Dynamical friction and viscous stirring increase particle velocities in models with large (8 km radius) initial bodies. This velocity increase delays runaway growth by a factor of two compared to models without velocity evolution.  In contrast, collisional damping dominates over dynamical friction and viscous stirring in models with small (80--800 m radius) initial bodies.  Collisional damping decreases the time scale to runaway growth by factors of 4--10 relative to constant velocity calculations.  Simulations with minimum mass solar nebulae, $M_0 \\sim$ 10 \\mearth, and small eccentricities, $e \\approx 10^{-3}$, reach runaway growth on time scales of 20--40 Myr with 80 m initial bodies, 50--100 Myr with 800 m bodies, and 75--250 Myr for 8 km initial bodies.  These growth times vary linearly with the mass of the annulus, $\\tau_r \\propto M_0^{-1}$, but are less sensitive to the initial eccentricity than constant velocity models.  In both sets of models, the time scales to produce 1000+ km objects are comparable to estimated formation time scales for Neptune. Thus, Pluto-sized objects can form in the outer solar system in parallel with the condensation of the outermost large planets. ",
        "introduction": "Current models for planet formation involve aggregation of solid planetesimals and gas accretion in a circumstellar disk (see, for example, \\cite{hay85}; \\cite{bos93}; and references therein). Large dust grains within the disk first settle to the midplane. These grains may then coagulate into successively larger grains (e.g., \\cite{wei80}; \\cite{wei93}) or continue to settle in a very thin layer that eventually becomes gravitationally unstable (e.g., \\cite{gol73}). Both paths produce km-sized planetesimals that collide and merge to produce large bodies such as planets.  Despite the complex and sometimes unknown physics, many simulations produce objects resembling known planets on time scales roughly comparable to the expected lifetime of the protosolar nebula (see, for example, \\cite{saf69}; \\cite{gre78}, 1984; \\cite{nak83}; \\cite{wet89}; 1993; \\cite{spa91}; \\cite{kol92}; \\cite{wei92}; \\cite{pol96}).  Recent observations of slow-moving objects in the outer solar system offer a new challenge to planet formation models.  The trans-Neptunian region now contains several dozen Kuiper Belt objects (KBOs) with estimated radii of 100--300 km (\\cite{jew96}).  The orbits of known KBOs suggest a division into at least three dynamical components with an inner radius of 30 AU and an unknown outer radius: (i) the classical KBOs, objects with roughly circular orbits, (ii) the resonant KBOs, objects in orbital resonance with Neptune (\\cite{jew96}), and (iii) the scattered KBOs, objects with large, eccentric orbits (\\cite{luu97}). Although the known population is still small, Jewitt \\etal (1996) estimate that the region between 30 and 50 AU contains $\\approx$ 70,000 objects larger than 100 km.  The total mass in the classical Kuiper Belt is thus at least 0.1 \\mearth.  This mass probably represents a small fraction of the initial mass, because dynamical interactions with Neptune reduce the number of KBOs on short time scales compared to the age of the solar system (\\cite{lev93}; \\cite{mal96}).  Despite these new observations, the origin of KBOs remains uncertain. Edgeworth (1949) and Kuiper (1951) first suggested that the Kuiper Belt was a natural extension of the original solar nebula. Holman \\& Wisdom (1993) later showed that small KBOs, once formed, can survive at 30--50 AU for times approaching the age of the solar system.  More recent dynamical studies confirm this conclusion and explain the observed distribution of KBOs in a general way (\\cite{lev93}).  The formation process and time scale for KBOs, however, is still controversial.  Planetesimal simulations for plausible protosolar nebulae at 25--30 AU show that Neptune can grow to its present size in 10--100 Myr (\\cite{fer81}, 1984; \\cite{ip89}; \\cite{pol96}).  These results suggest that Pluto might form on a similar time scale at $\\sim$ 40 AU, because growth times are not a steep function of semi-major axis (see \\cite{aar93}; \\cite{pol96}; and references therein).  Nevertheless, Stern (1995, 1996) and Stern \\& Colwell (1997) conclude that KBO formation requires 100--1000 Myr for the conditions expected in the outer solar system. These limits far exceed the time scale required to produce Neptune, whose accretion time is constrained by the 10--100 Myr lifetime of the protosolar nebula (see \\cite{pol96} and references therein).  In this paper, we attempt to resolve the uncertainties surrounding KBO production with a new planetesimal simulation at 35 AU.  We have developed an evolution code to follow the growth and velocity evolution of planetesimals with a wide range of initial masses.  The code matches analytic models and duplicates Wetherill \\& Stewart's (1993; hereafter WS93) simulation of planetesimal evolution at 1 AU.  Our numerical results demonstrate that small bodies with initial radii of 80 m to 8 km can produce 1000+ km objects on time scales of 10--100 Myr.  We confirm these calculations with a simple analytic estimate of the growth time as a function of semi-major axis.  This analysis supports previous estimates for a short growth phase for Neptune, 10--100 Myr, and indicates that Pluto-Charon can form just outside the current orbit of Neptune on a similar time scale.  We outline the accretion model in Sec. 2, describe our calculations in Sec. 3, and conclude with a discussion and summary in Sec. 4. The Appendix contains a complete description of the algorithms and detailed comparisons with analytic models. ",
        "conclusions": "We have developed a time-dependent planetesimal evolution program similar to the WS93 code used to simulate the formation of terrestrial embryos from small bodies. The program incorporates coagulation with realistic cross-sections and velocity evolution using the statistical formulation of Barge \\& Pellat (1990, 1991; see also Hornung \\etal 1985).  Our numerical solutions to the coagulation equation agree with analytical solutions for three standard test cases.  Our results also agree with WS93's simulation of the formation of the Earth at 1 AU.  The present models do not incorporate fragmentation of bodies during collisions. We will include fragmentation in a separate paper.  We have considered two simple cases of planetesimal evolution in a 6 AU wide annulus centered at 35 AU.  Models without velocity evolution invariably produce several large bodies that accrete practically all of the material in the annulus. The runaway growth in these simulations occurs without dynamical friction or gas drag; it is a direct consequence of gravitational focusing.  The time required to produce a runaway body in our models scales inversely with the initial mass of the annulus {\\it and} with the initial radii and velocities of the planetesimals. For bodies with $r_0$ = 80--8000 m and $e_{init} = 10^{-3}$, our simulations produce runaway growth in 100 Myr for annular masses of roughly 10--30 $M_{\\earth}$.   The time scale for runaway growth increases to 700--2000 Myr for $e = 10^{-2}$. A minimum mass solar nebula with $\\Sigma \\propto R^{-3/2}$ contains 7--15 \\mearth~in a 6 AU wide annulus centered at 35 AU.  These models thus reach runaway growth in a minimum mass nebula on time scales comparable to the maximum formation time scale for Neptune ($\\sim$ 50-100 Myr, \\cite{lis96}; \\cite{pol96}). Runaway growth on much shorter time scales, $\\sim$ 10 Myr, requires annular masses that far exceed the minimum mass solar nebula, $\\sim$ 100 $M_{\\earth}$ in a 6 AU wide annulus for $e_{init} = 10^{-2}$ to $10^{-3}$.  Models with velocity evolution produce runaway growth on a much wider range of time scales compared to constant velocity calculations. First, dynamical friction and viscous stirring dominate the evolution of models with relatively large (8 km) initial bodies.  The velocities of these bodies thus increase as collisions produce more massive objects. This velocity increase delays runaway growth by factors of 2 or more compared to constant velocity evolution.  The delay in the runaway growth time increases with increasing $e_{init}$. In contrast, collisional damping dominates the evolution of models with smaller (80--800 m) initial bodies.  These bodies `cool' until the largest objects have radii of 10--20 km.  Dynamical friction and viscous stirring then `heat up' the small bodies but this heating is small compared to the velocity increases of the $r_0$ = 8 km models. For $e_{init}$ = $10^{-3}$, collisional damping enhances collision rates and decreases the time scale to runaway growth by factors of 4--12 compared to constant velocity calculations.  Our simulations of minimum mass solar nebulae with $e_{init} = 10^{-3}$ reach runaway growth on time scales of 20--40 Myr for 80 m initial bodies, 50--100 Myr for 800 m bodies, and 75--250 Myr for 8 km bodies. These time scales increase by factors of 2--4 for $e_{init} = 10^{-2}$.  The formation of runaway bodies in constant velocity simulations is surprising.  Previous simulations demonstrate a need for dynamical friction, which decreases the velocities of the most massive objects (\\cite{wet89}, 1993; \\cite{bar90}, 1991; \\cite{oht92}; \\cite{id92a}, 1992b; \\cite{kok96}).  Gravitational focusing then allows these bodies to sweep up lower mass bodies very rapidly. Kokubo \\& Ida (1996) summarize necessary and sufficient conditions for runaway growth and show that the ratio of the maximum mass of planetesimals to their mean mass increases dramatically during runaway growth.  Our constant velocity simulations satisfy the ``necessary'' condition for runaway growth,  \\begin{equation} 2 R_{H,ij} \\Omega \\simless V_{ij} \\simless V_{e,ij} \\end{equation}  \\noindent but do not meet the ``sufficient'' condition,  \\begin{equation} \\frac{dV}{dm} < 0 ~ , \\end{equation}  \\noindent because our velocities are constant with mass, $dV/dm$ = 0. Nevertheless, the maximum mass, $M_{max}$, of each constant velocity simulation increases much more rapidly than the mean mass, $<m_i>$ (Figure 14).  In Figure 14(a), the ratio $M_{max} / <m_i>$ increases slowly during the orderly growth phase and then increases rapidly when several isolated bodies begin to accrete most of the mass in the annulus.  This runaway growth is not as extreme as that seen in models with velocity evolution.  We derive $M_{max} / <m_i> ~ \\approx 10^6$ in constant velocity models compared to $M_{max} / <m_i> ~ \\approx 10^8$--$10^{10}$ in models with velocity evolution (Figure 14(b)).  Dynamical friction is responsible for the larger increase in $M_{max} / <m_i>$ in models with velocity evolution.  Although the constant velocity simulations are artificial, they are a useful guide for planetesimal growth in the outer solar system. Our results show that collisional damping by small bodies can overcome viscous stirring and keep particle velocities low at $a$ = 35 AU.  The relative efficiency of collisional damping should increase with $a$, because the collision rates decrease more rapidly with $a$ than do the damping rates ($\\tau_{coll} \\propto a^{-5/2}$ vs. $\\tau_{damp} \\propto a^{-3/2}$).  The situation in the outer solar system differs markedly from conditions at small $a$. In our simulations at 1 AU, bodies with $r_i \\approx$ 80--800 m grow to 10 km objects in $\\sim$ 1000 yr.  As in the 35 AU simulations described above, collisional damping reduces the velocities of the small bodies by a factor of $\\sim$ 2 in 1000 yr.  In contrast to our 35 AU models, viscous stirring and dynamical friction act quickly to increase velocities once larger bodies are produced at 1 AU.  Aside from the longer time scale to reach runaway growth, this evolution then follows closely the simulations in Figures 1--2.  With these considerations in mind, we suggest a modest modification to the sufficient condition for runaway growth,  \\begin{equation} \\frac{dV}{dm} \\le 0 ~ , \\end{equation}  \\noindent instead of the condition in Eqn. (3).  This condition maintains the relative growth rate necessary for runaway in a three-dimensional system (see \\cite{kok96})  \\begin{equation} \\frac{1}{M} \\frac{dm}{dt} \\propto M^{1/3} ~ , \\end{equation}  \\noindent when the gravitational focusing factor is large (e.g., $V_{ij} \\simless V_{e,ij}$).  In addition to this new criterion for runaway growth, our simulations show that Neptune and Pluto can grow in parallel at $a \\approx$ 30--35 AU.  For a minimum mass solar nebula, previous calculations indicate that Neptune reaches its current size in no longer than 50--100 Myr (\\cite{ip89}; \\cite{lis96}; \\cite{pol96}).  In our models of minimum mass solar nebulae at 35 AU, 800 m objects grow to Pluto-sized planets on similar time scales.  However, recent observations indicate a much shorter lifetime, $\\sim$ 5--10 Myr, for gaseous disks surrounding nearby pre-main sequence stars (\\cite{sar93}; \\cite{str93}; \\cite{rus96}; see also \\cite{fer97}).  If our solar system evolved on a similar time scale, the formation of the gas-rich outer planets requires solar nebulae with masses 2--5 times larger than the minimum mass solar nebula, $M_{min}$ (\\cite{lis96}; \\cite{pol96}).  Neptune then attains its current size in 5--10 Myr.  At 35 AU, objects reach 1000 km sizes in 10--20 Myr for $M_0 \\approx$ 2--3 $M_{min}$ and in 5--10 Myr for $M_0 \\approx$ 5 $M_{min}$ assuming the initial bodies have small masses and eccentricities. Fragmentation should not change these conclusions unless collisional erosion can prevent the formation of 1 km bodies from smaller building blocks.  Our Kuiper Belt simulation produce Pluto-sized objects on reasonable time scales for other plausible solar nebula models. In particular, Cameron's (1995 and references therein) detailed disk models for the protosolar nebula have a shallower density density distribution, $\\Sigma = \\Sigma_0 (a/a_0)^{-1}$, than the minimum mass model we considered above to derive our success criteria (see also \\cite{rud86}; \\cite{rud91}). Cameron's model contains $\\sim$ 100 \\mearth~in our 6 AU annulus. The growth time for 1000+ km objects scales simply from results with smaller disk masses.  For models with velocity evolution, we estimate $\\tau_r \\approx 3$ Myr for $r_0$ = 80 m, $\\tau_r \\approx 8$ Myr for $r_0$ = 800 m, and $\\tau_r \\approx 15$ Myr for $r_0$ = 8 km. Although all of these models can produce Pluto-sized objects before Neptune reaches its final mass, models with $r_0$ = 8 km produce more KBOs with $r_i \\sim$ 100--300 km as outlined in Sec. 3.3.  These results -- together with recent dynamical calculations of Malhotra (1993, 1995, 1996) -- suggest a self-consistent picture for the formation of Pluto-Charon in the outer solar system. In this picture,  Pluto and Charon begin as $\\sim$ 1 km planetesimals at $a \\sim$ 35--40 AU and grow to their present sizes on a time scale of 10--100 Myr.  Both objects are runaway bodies more massive than the bulk of the planetesimal mass distribution and have low velocities due to collisional damping and dynamical friction.  At a somewhat smaller semi-major axis, $a \\approx$ 25--30 AU, Neptune accretes its current mass in 10--100 Myr and migrates radially outward through the protosolar disk during the late stages of giant planet formation (\\cite{mal93}; \\cite{pol96}; \\cite{fer96}; see also \\cite{ipa89}, 1991). During this outward migration, Neptune captures Pluto-Charon in the 3:2 resonance (\\cite{mal95}).  This capture should effectively end further growth of Pluto-Charon, because the orbital elements increase to $e \\sim$ 0.2 and $i \\sim 10^{\\circ}$ on a short time scale, $\\sim$ 10 Myr, inside the resonance (\\cite{mal95}). High velocity collisions within the resonance should also hinder growth as in our large $e$ models.  Neptune also captures other KBOs at $a \\approx$ 35--40 AU into the 3:2 and other resonances (\\cite{mal96}; see also \\cite{jew96}). Further growth of these objects is also slowed due to rapidly increasing velocities (\\cite{mal95}, 1996; \\cite{mor95}; \\cite{dav97}). As long as collisional erosion does not decrease significantly the radii of captured KBOs, this sequence of events accounts for the general aspects of the mass distribution and orbital elements of observed KBOs in a simple way.  This picture for the formation of Neptune, Pluto-Charon, and KBOs differs from those of Stern (1995, 1996) and Stern \\& Colwell (1997a,b), who studied KBO evolution at 40--70 AU.  Stern \\& Colwell (1997) concluded that the formation of Neptune and KBOs, including Pluto-Charon, requires $\\sim$ 100--1000 Myr due to the low collision rates at $a \\approx$ 30--50 AU.  Stern \\& Colwell's (1997a; also \\cite{ste95}, 1996) time scale for KBO formation is a factor of 10--20 longer than our time scale.  Their results also conflict with the more detailed gas dynamic calculations of Pollack \\etal (1996). Although the exact origin of this discrepancy is unclear, we suspect mass spacing, collision rates, and (less probably) fragmentation may be responsible. Our mass spacing, $\\delta$ = 1.4, should produce more accurate estimates for the growth time than $\\delta$ = 2 (\\cite{st97a}) or $\\delta$ = 4 (\\cite{ste96}).  We estimate 10\\%--20\\% delays in runaway growth for $\\delta$ = 2 and expect very long delays for $\\delta$ = 4.  In addition, our collision rates include a factor, $\\beta_{coll}$, that accounts for the gaussian distribution of impact velocities (\\cite{grz92}; see Eqn. A14).  Neglecting this factor delays runaway growth by a factor of 3 (see also \\cite{ws93}). Stern (1995) does not include this factor in his cross-section (see his Eqn. 5) and thus derives much longer growth times.  As noted above, our neglect of fragmentation encourages runaway growth.  We do not think that including this process should delay runaway growth by another factor of 3, for reasons outlined earlier, and plan to test this suspicion in our next paper.  However the theoretical issues may be resolved, several consequences of our accretion models can be tested with additional observations. All models that reach runaway growth produce 10--100 Pluto-sized objects with radii $r_{max} \\approx$ 1000--2000 km and a roughly power-law mass distribution with a maximum radius at $\\sim$ 0.1--0.2$r_{max}$ (see Figures 8--10).  If Pluto and Charon are runaway bodies produced at $a \\approx$ 35--40 AU, there should be several additional ``Plutos'' with similar orbital elements.  Malhotra (1995) reached a similar conclusion and noted that recent searches have not excluded the possibility of several additional Plutos in 3:2 orbital resonance with Neptune.  The accretion models also predict a cumulative power law distribution, $N_C \\propto r^{-3}$, for objects with $r \\lesssim$ 100--200 km.  By analogy with WS93, this shape is fairly independent of fragmentation as long as collisions produce overall growth instead of disruption.  The best fit to the observed distribution is shallower than expected, $N_C \\propto r^{-2}$, but the data are not yet accurate enough to preclude our model prediction (\\cite{jew96}). Larger surveys will provide a better test of this prediction (\\cite{luu97}).  Finally, our results suggest that KBO formation is likely in other solar systems.  KBOs can grow in the dusty disks that surround many nearby main-sequence stars, if the disk masses are within an order of magnitude of the ``maximum'' disk masses for $\\alpha$ Lyr, $\\alpha$ PsA, and $\\beta$ Pic (\\cite{bac93}).  Stars with disk masses near the minimum dust masses of these well-studied A-type stars probably have few, if any, KBOs, but could have many objects with $r_i \\sim$ 1--10 km.  Observations of nearby pre-main sequence stars also indicate substantial masses, $M_d \\gtrsim 10^3~\\mearth$, in circumstellar disks with radii of 100--1000 AU (e.g., \\cite{sar93}; see also \\cite{bec96}).  These data imply $M_0$ $\\sim$ 10--100 \\mearth~in the Kuiper Belt.  With formation time scales of 1--10 Myr, KBOs can grow in massive circumstellar disks during the pre-main sequence contraction phase of a low mass star (see \\cite{kh95} and references therein).  In a less massive disk,  KBOs grow while the central star is on the main sequence.  We expect no KBO formation in circumstellar disks with very low masses, $M_d \\lesssim$ 10 \\mearth. These disks can produce 1--10 km objects at $a \\gtrsim$ 30 AU, unless fragmentation prevents growth of icy bodies.  We will explore this possibility in a subsequent paper.  \\vskip 6ex  We thank B. Bromley for making it possible to run our code on the JPL Cray T3D and for a generous allotment of computer time through funding from the NASA Offices of Mission to Planet Earth, Aeronautics, and Space Science.  K. Kirby and J. Babb also graciously allocated us computing time on the CfA Dec Alpha computer.  Comments from A. Cameron, F. Franklin, M. Geller, M. Holman, J. Wood and an anonymous referee greatly improved our presentation.  We acknowledge G. Stewart for clarifying many details of the WS93 calculations.  We also thank S. Stewart for designing and implementing a preliminary version of our computer code as part of her Harvard Senior Thesis.  Finally, S.K. thanks S. Starrfield for teaching the tools for solving time-dependent differential equations.  \\appendix"
    },
    "9804/astro-ph9804095_arXiv.txt": {
        "abstract": "New BVI CCD photometry is presented for 60 RR Lyrae variables in the globular cluster M3. Light curves have been constructed and ephemerides have been (re)-derived for all of them. Four stars (i.e. V29, V136, V155 and V209), although recognized as variables, had no previous period determinations. Also, the period derived for V129 is significantly different from the one published by Sawyer-Hogg (1973). Light curve parameters, i.e. mean magnitudes, amplitudes and rise-times, have been derived.  The discussion of these results in the framework of the stellar evolution and pulsation theories will be presented in a forthcoming paper. ",
        "introduction": "Multi-colour observations of RR Lyrae variables are extremely important to obtain information on their physical parameters, such as temperature, mass, and luminosity, which are essential to the understanding of the different characteristics of these stars in different clusters. One of the most significant differences is related to the separation of globular clusters into two groups according to the mean period of the RRab variables, which was found by Oosterhoff (1939, 1944), namely $< P_{ab} >$=0.55 days for group I (OoI) and $< P_{ab} >$=0.65 days for group II (OoII). This ``dichotomy'', later found to be related to the cluster metallicity (Arp 1955), has been the basis for many subsequent studies which have led several authors, {\\it in primis} Sandage (1981, 1990, 1993 and references therein) to define a luminosity-metallicity relation among RR Lyrae stars as a result of this observed ``period-shift''. The fiducial cluster in the study of this period-shift effect has traditionally been M3 for the OoI group. The data for the M3 RR Lyraes, however, were still essentially the UBV photographic photometry by Roberts and Sandage (1955), Baker and Baker (1956) and Sandage (1959), therefore a program was started in order to obtain more accurate CCD photometry for the study of the characteristics of both the Colour-Magnitude Diagram and the RR Lyrae variables in this cluster.  M3 (NGC5272; RA=13 42 11, DEC=+28 22 32 (2000); l=42$\\deg$ b=79$\\deg$) is among the most important and prominent clusters in the northern hemisphere. It is located at $\\sim$ 11.6 kpc from the Galactic Center and 9.5 kpc above the plane, and is at 9.7 kpc from the Sun (Harris 1996). It contains the largest number of variable stars known within a single cluster (260 are quoted by Suntzeff, Kinman \\& Kraft 1991) with a rather high specific frequency of RR Lyraes (i.e. normalized to total mass). The frequency of double-mode pulsators, however, is quite low: only three RRd are known (Nemec \\& Clement 1989, hereafter NC89, and Kaluzny et al. 1997, hereafter K97), two of which (V68 and V87) have been known to be RRd stars for some time, and the third one (V79) has been identified as an RRd star only recently because it appears to have switched modes some time between 1962 and 1996 (Clement et al. 1997). Its metallicity, for a long time considered to be [Fe/H] $\\sim$ -1.6, has been recently revised using high resolution spectroscopic data and appears to be somewhat higher, i.e. $\\sim$ -1.4 (Carretta \\& Gratton 1997).  M3 has been the target of a very large number of studies and analyses. For what concerns the photometry of RR Lyrae variables, we mainly refer to Roberts \\& Sandage (1955), Baker \\& Baker (1956) and Sandage (1959), which constituted the largest body of photometric data and was used as a basis for all subsequent studies of the global properties of RR Lyraes. Several more references on individual stars are listed in Sawyer-Hogg (1973 - hereafter SH73) and in the updated compilation of the SH73 Catalogue by Clement (1996). The most recent CCD observations of RR Lyraes in M3 are from K97 and from M. Corwin (private communication).  In this paper we present a new photometric study of 60 RR Lyrae variable stars in M3, based on BVI CCD observations taken over a period of 5 years and a new independent absolute photometric calibration (described in detail by Ferraro et al. 1997, hereafter F97). Here we will present the observations and an initial analysis of the data. A more detailed analysis of the RR Lyrae properties is deferred to a separate paper (Cacciari et al. 1998), whereas the new photometry of the extended sample of non-variable stars ($\\sim$ 19,000 stars) and related CMD have been presented and discussed by Buonanno et al. (1994: hereinafter B94) and F97. ",
        "conclusions": ""
    },
    "9804/astro-ph9804160_arXiv.txt": {
        "abstract": "We present the first intermediate-resolution ($\\lambda$ / 3000) spectrum of the bright quasi-stellar object 3C273 at wavelengths between 900 and 1200 \\AA. Observations were performed with the Berkeley spectrograph aboard the \\orf-II mission \\cite{Hetal98}. We detect Lyman $\\beta$ counterparts to intergalactic Lyman $\\alpha$ features identified by \\citeN{MWSG91} at $cz$ = 19900, 1600, and 1000 \\kms; counterparts to other putative Lyman $\\alpha$ clouds along the sight line are below our detection limit. The strengths of the two very low redshift Lyman $\\beta$ features, which are believed to arise in Virgo intracluster gas, exceed preflight expectations \\cite{WRWMH95}, suggesting that the previous determination of the cloud parameters may underestimate the true column densities. A curve-of-growth analysis sets a minimum \\hxi\\ column density of $4 \\times 10^{14}$ cm$^{-2}$ for the 1600 \\kms\\ cloud. We find marginally significant evidence for Galactic \\htwo\\ along the sight line, with a total column density of about $10^{15}$ cm$^{-2}$. We detect the stronger interstellar \\oxvi\\ doublet member unambiguously; the weaker member is blended with other features. If the Doppler $b$ value for \\oxvi\\ is comparable to that determined for \\nxv\\ \\cite{SST97} then the \\oxvi\\ column density is $7 \\pm 2 \\times 10^{14}$ cm$^{-2}$, significantly above the only previous estimate \\cite{D93}. The \\oxvi / \\nxv\\ ratio is about 10, consistent with the low end of the range observed in the disk (compilation of \\citeNP{HB96}). Additional interstellar species detected for the first time toward 3C273 (at modest statistical significance) include \\pxii, \\fexiii, \\arxi, and \\sxiii. ",
        "introduction": "We present the first intermediate-resolution far-ultraviolet spectrum of the bright quasi-stellar object 3C273 in the 900 -- 1210 \\AA\\ band, emphasizing absorption features arising along the line of sight. \\citeN{Appetal98} discuss the intrinsic spectrum of 3C273. These data were collected with the Berkeley spectrograph \\cite{Hetal98} during the ORFEUS-SPAS II mission. The \\orf\\ project and the {\\it ASTRO-SPAS\\/} platform are described in \\citeN{Getal91}.  3C273 was observed five times, for a total of 10,797 s. All observations took place in late 1996, between 330/21:41 and 338/16:33 (Day of Year/HH:MM, GMT). The target coordinates were $\\alpha$ = 12 29 06.7, $\\delta$ = +02 03 09 (J2000). Absolute positioning of the 26$\\arcsec$ diameter ORFEUS entrance aperture is accurate to $\\pm$ 5$\\arcsec$. Extraction of the spectrum and subtraction of airglow follow the discussion in \\citeN{Hetal98}. Statistical and systematic uncertainties associated with background and airglow subtraction are properly tracked. As a final step, we bin the data on 0.1 \\AA\\ centers.  In Figure 1 we show the complete spectrum (solid line), the 1 $\\sigma$ uncertainty associated with shot noise and detector flat-field effects (dotted line), and a continuum established by an automated fitting routine (dashed line).  The automated fitting routine identifies absorption features deeper than a statistically determined threshold, replaces those data with a linear interpolation from nearby wavelengths, smoothes heavily, then iterates twice more at increasing sensitivity to absorption lines. The resulting continuum has the advantage of being determined objectively, and evidently tracks most of the true spectrum well at wavelengths greater than 1000 \\AA. Like many such routines, however, this one yields unreliable results when strong features are closely spaced on a curved continuum (as in the region of absorption lines 3 and 4) and will tend to underestimate the true value when the spectrum contains a cluster of weak absorption features below the statistical threshold (note the weak dip in the fitted continuum around 1050 \\AA).  In Table 1 we list the interstellar absorption features and blends that are sufficiently strong and isolated for direct equivalent width measurement using the continuum shown in Figure 1 (unless otherwise noted). Statistical noise in the spectrum (1$\\sigma$) corresponds to an unresolved absorption feature with equivalent width of about 0.040 \\AA; continuum placement uncertainty corresponds to about half that value. The value in the fourth column includes the statistical uncertainty only. For species whose column densities have not previously been reported, we list the minimum column (if optically thin) based on the best fit equivalent width and the oscillator strenths of \\citeN{M91}.  \\begin{deluxetable}{llllll} \\tablewidth{0pt} \\tablecaption{Interstellar Lines and Equivalent Widths}  \\tablehead{ \\colhead{\\#}            & \\colhead{W$_{obs}$ (\\AA)}           & \\colhead{W$_{eq}$ (\\AA)}           & \\colhead{$\\sigma$ (\\AA)}      & \\colhead{ID}        & \\colhead{Log ($N_{min}$) (cm$^{-2}$)} } \\startdata  1 & 1206.44 & 0.573 & 0.056 & Si III 1206.50 & 13.4 \\nl 2 & 1199.75 & 0.896 & 0.114 & N I blend & \\nl 3\\tablenotemark{a} & 1193.56 & 0.510 & 0.050 & Si II 1193.29 & \\nl 4\\tablenotemark{a} & 1190.58 & 0.532 & 0.051 & Si II 1190.42 & \\nl 5 & 1152.91 & 0.105 & 0.038 & P II 1152.82 & 13.6 \\nl 6 & 1144.92 & 0.340 & 0.054 & Fe II 1144.94 & \\nl 7 & 1143.44 & 0.137 & 0.040 & Fe II 1143.22 & \\nl 8 & 1134.63 & 1.135 & 0.084 & N I blend & 15.1 \\nl 9 & 1122.54 & 0.157 & 0.038 & Fe III 1122.53 & 14.3 \\nl 10 & 1121.99 & 0.149 & 0.038 & Fe II 1121.97 & \\nl 11 & 1096.94 & 0.202 & 0.047 & Fe II 1096.88 & \\nl 12 & 1084.19 & 0.515 & 0.051 & N II 1083.99, N II* blend & \\nl 13\\tablenotemark{b} & 1066.77 & 0.135 & 0.047 & Ar I 1066.66 & \\nl 14 & 1063.20 & 0.327 & 0.049 & Fe II 1063.18 + \\htwo & \\nl 15\\tablenotemark{c} & 1048.19 & 0.096 & 0.035 & Ar I 1048.22 & 13.6 \\nl 16 & 1039.18 & 0.433 & 0.044 & O I 1039.23 & \\nl 17 & 1031.87 & 0.363 & 0.038 & O VI 1031.93 & \\nl 18 & 1020.75 & 0.129 & 0.040 & Si II 1020.70 & \\nl 19 & 1012.53 & 0.083 & 0.036 & S III 1012.50 & 14.4 \\nl 20 & 1008.51 & 0.165 & 0.039 & NOID & \\enddata  \\tablenotetext{a}{Continuum determined with 2nd order local polynomial fit.} \\tablenotetext{b}{Up to 0.02 \\AA\\ may be contributed by \\htwo\\ ($J$~=~2).} \\tablenotetext{c}{Continuum may be depressed by adjacent weak lines.}  \\end{deluxetable}  The identification of most of the strong interstellar features listed in Table 1 is secure \\cite{M91}. There are no obvious interstellar or intergalactic candidates for the unidentified line at 1008.5 \\AA, which may be spurious. Although the observed features at 1066.8 and 1048.2 \\AA\\ presumably correspond to a pair of \\arxi\\ lines, their equivalent widths at first glance appear to be anomalous (the 1048.2 \\AA\\ line should be stronger). As noted, however, the longer wavelength feature may be partially blended with an \\htwo\\ line at 1066.9 \\AA. The shorter wavelength feature probably suffers from local continuum depression by nearby \\htwo. Using the linear continuum shown in Figure 2, we estimate a best fit equivalent width closer to 0.13 \\AA. These systematic effects, in combination with the statistical errors, do much to reconcile the anomaly (especially if the features are on the flat part of the curve-of-growth).  Below 1000 \\AA, the increasing noise in the spectrum makes quantitative analysis of absorption features difficult, and the continuum fitting routine fails. Galactic features that are clearly detected include a blend from \\oxi\\ (and \\nxiii ?) near 989 \\AA, \\cxiii\\ (977.02 \\AA), Lyman $\\delta$ (972.54 \\AA), and an \\nxi\\ blend near 953.8 \\AA. Higher series Lyman lines are likely to be filled in by diffuse emission (we correct for diffuse \\hxi\\ emission only through Lyman $\\delta$). An intriguing dip near 982 \\AA\\ may be associated with \\cxiii\\ in the Virgo cluster, but this identification is speculative. Difficulties in continuum placement render the equivalent widths very uncertain; we attempt no quantitative analysis of absorption lines below 1000 \\AA. ",
        "conclusions": "\\label{conclusions}  We observed the bright quasi-stellar object 3C273 for 10,797 s with the Berkeley spectrograph aboard the \\orf-II mission \\cite{Hetal98}. The resulting spectrum offers intermediate spectral resolution ($\\lambda$ / 3000 FWHM) and a 1$\\sigma$ noise corresponding to an absorption feature with equivalent width of about 0.040 \\AA.  The spectrum reveals, at modest statistical significance, a variety of previously undetected low ionization interstellar species, including \\pxii, \\fexiii, \\arxi, and \\sxiii. There is marginally significant evidence for \\htwo, with a total column density of about $10^{15}$ cm$^{-2}$. The $J$~=~0 and $J$~=~1 rotational levels seem comparably populated, consistent with an excitation temperature $T_{01}$ typical for interstellar clouds (50 -- 80 K).  We detect the stronger interstellar \\oxvi\\ doublet member unambiguously; the weaker member is blended. If the Doppler $b$ value for \\oxvi\\ is comparable to that determined for \\nxv\\ \\cite{SST97}, the \\oxvi\\ column density is $7 \\pm 2 \\times 10^{14}$ cm$^{-2}$. The \\oxvi / \\nxv\\ ratio is about 10, consistent with the low end of the range observed in the disk (compilation of \\citeNP{HB96}). If impulsively heated gas cools as it moves away from the Galactic disk, the \\oxvi / \\nxv\\ ratio might be expected to decrease (compared to the disk value) along halo sight lines.  We detect Lyman $\\beta$ counterparts to intergalactic Lyman $\\alpha$ features identified by \\citeN{MWSG91} at $cz$ = 19900, 1600, and 1000 \\kms. Counterparts to other putative Lyman $\\alpha$ clouds along the sight line are below our detection limit. The strengths of the two very low redshift Lyman $\\beta$ features, which are believed to arise in Virgo intracluster gas, exceed preflight expectations \\cite{WRWMH95}. A curve-of-growth analysis sets a minimum \\hxi\\ column density of $4 \\times 10^{14}$ cm$^{-2}$ for the 1600 \\kms\\ cloud. Following \\citeN{GBJS93}, our revised column density and nondetection of 21 cm radio emission discussed in that work can set a lower limit of about 6 kpc to the size of the Virgo cluster clouds."
    },
    "9804/astro-ph9804026_arXiv.txt": {
        "abstract": "s{ The X-ray emission from clusters of galaxies is one of the best observational probe to investigate the distribution of dark matter at intermediate and high redshifts.  Since the disposition of the intracluster plasma (ICP) responsible of the emission is crucial to link X-ray properties to the global properties of the dark matter halos, we propose a semi--analytical approach for the diffuse baryons. This comprises the following blocks: Monte Carlo ``merging histories'' to describe the dynamics of dark matter halos; the central {\\sl hydrostatic} disposition for the ICP; conditions of shock, or of closely adiabatic compression at the {\\sl boundary} with the external gas, preheated by stellar energy feedbacks. From our model we predict the $L-T$ correlation, consistent with the data as for shape and scatter.  }  \\normalsize\\baselineskip=15pt ",
        "introduction": "Groups and clusters of galaxies constitute cosmic structures sufficiently close to equilibrium and with sufficient density contrast ($\\delta\\approx 2\\, 10^2$ inside the virial radius $R$) as to yield definite observables. They are dominated by dark matter (hereafter DM), while the baryon fraction is observed to be less than $20$\\%.  The great majority of these baryons are in the form of {\\sl diffuse plasma} (ICP) with densities $n\\sim 10^{-3}$ cm$^{-3}$ and virial temperatures $k\\,T\\sim GM\\,m_H/10 R\\sim 5$ keV, and are responsible for powerful X--ray luminosities $L\\sim 10^{44}$ erg/s by optically thin thermal bremsstrahlung.  As the plasma is a good tracer of the potential wells, much better than member galaxies, the X--ray emission is a powerful tool to investigate the mass distribution out to moderate and high redshifts.  The ICP temperature directly probes the height of the potential well, with the baryons in the  role of  mere tracers; on the other hand, the luminosity, with its strong dependence on density ($L\\propto n^2$), reliably probes the baryonic content and distribution. Statistically, a definite $L$--$T$ correlation is observed (albeit with considerable scatter), and this provides the crucial link to relate the X-ray luminosity functions with the statistics of the dark mass $M$ or with that of the corresponding $T$.  Many numerical experiments, using hydrodynamical N--body (see especially Gheller, Pantano \\& Moscardini 1997, Bryan \\& Norman 1997), provide a comprehensive tool to model the ICP distribution in the potential wells.  However, such numerical experiments still do not have enough dynamic range to describe DM and ICP over the full range from $\\sim 50$ Mpc associated with the large scale structures (which guide the ongoing mergers of DM halos), to the inner $50$ kpc where the ICP yields a substantial contribution to $L$, and moreover do not include properly non--gravitational effects, which instead cannot be ignored.  The simplest semi--analytical approach is constituted by the Self Similar model (Kaiser 1986) which include only gravity and assumes the ICP amount to be proportional to the DM at all $z$ and $M$.  This leads to a relation $L\\propto T^2$ conflicting with the observed correlation for rich clusters, which is close to $L\\propto T^3$ (David et al. 1993; Mushotzky \\& Scharf 1997).  A missing ingredient is thought to be the stellar energy feedback by supernovae.  The above features motivate us to develop a comprehensive semi--analytical model for the ICP which include the effects of stellar energy scales.  First we assume that such energy input is efficient in depleting the potential wells of the clusters progenitors, at $z\\simeq 1 \\div 2$, and in pre-heating the intergalactic medium to temperatures in the range $T_1=0.1 \\div 0.8$ keV as recently evidentiated in the outer cluster atmosphere (Henriksen \\& White 1996).  Then we describe clusters evolution as a sequence of hierarchical merging episodes of the DM halos, associated in the ICP to shocks of various strengths (depending on the mass ratio of the merging clumps), which provide the boundary conditions for the ICP to re--adjust to a new hydrostatic equilibrium.  In \\S 3 we shall show that our predictions for the X-ray properties of clusters are consistent with recent data (see also Cavaliere, Menci \\& Tozzi 1997, hereafter CMT97). ",
        "conclusions": "The ICP {\\it state} in the hierarchically evolving gravitational wells constitutes the focus of our new approach.  We propose that such state follows suit, passing through a sequence of equilibrium condition that we compute semi-analytically.  These computations comprise: the merging histories of the DM potential wells, obtained with a large statistics from Monte Carlo simulations of the hierarchical clustering; the inner {\\sl hydrostatic} equilibrium disposition, updated after each merging episode; and the {\\sl boundary} conditions provided by strong and weak shocks, or even by a closely adiabatic compression, depending on the ratio of the infall to the thermal energy in the preheated external medium.  The results of our model depend on two parameters, the external temperature $T_1$ and density $n_1$, which are not free. Specifically, we use for $T_1$ the range $0.1\\div 0.8 $ keV provided by the literature on stellar preheating.  The value of $n_1$ for rich clusters is related to the DM density by the universal baryonic fraction.  The expression of the bolometric luminosity is proportional to $g^2=(n_2/n_1)^2$, the square of the density jump at the bounding shock. The average of such factor over the merging histories coupled with $\\beta(T)$ is what gives to the statistical $L-T$ correlation the {\\sl curved} shape shown in fig. \\ref{fig4}b. In addition, our approach predicts an intrinsic variance of dynamical origin due to the different merging histories, and built in the factor $g^2$. Moreover the decreasing temperature profiles are in agreement with the published data and with the results from advanced simulations. A straightforward application is the prediction for the X-ray statistics in different FRW universes (see N. Menci, this volume, and Cavaliere, Menci \\& Tozzi 1998).      \\bigskip \\bigskip"
    },
    "9804/astro-ph9804210_arXiv.txt": {
        "abstract": "The X-ray colors and X-ray--to--optical luminosity ratios $(L_X/L_B)$ of 61 early-type galaxies observed with the {\\it ROSAT} PSPC are determined. The colors indicate that the X-ray spectral properties of galaxies vary as a function of $L_X/L_B$. The brightest X-ray galaxies have colors consistent with thermal emission from hot gas with roughly the same metallicity of 50\\% solar. The spatial variation of the colors indicates that the gas temperature in these galaxies increases radially. Galaxies with medium $L_X/L_B$ also have spectral properties consistent with emission from hot gas. If a simple one-component thermal model is assumed to describe the $0.1-2.0$ keV X-ray emission in these galaxies, then one possible explanation for the progressive decrease in $L_X/L_B$ among galaxies of this class could be the progressive decrease in metal abundance of the X-ray emitting gas contained by the galaxies. However, stellar X-ray emission may become a complicating factor for the fainter galaxies in this medium $L_X/L_B$ class.  Galaxies with the lowest $L_X/L_B$ values appear to be lacking a hot interstellar component. Their X-ray colors are consistent with those derived from the bulges of the spiral galaxies M31 and NGC~1291. In M31 the X-ray emission is resolved into discrete sources, and is apparently due primarily to low mass X-ray binaries (LMXBs). We therefore suggest that the bulk of the X-ray emission in the faintest ellipticals is also due to LMXBs. Previously, the X-ray spectra of X-ray faint galaxies had been found to be described by a hard component which was attributed to LMXB emission, and a very soft component of unknown origin. We show that the very soft component also likely results from LMXBs, as a very soft component is seen in the X-ray spectra of the nearby LMXB Her X-1 and LMXBs in the bulge of M31. If the X-ray emission in X-ray faint galaxies is primarily from stellar sources, then a range in $L_X/L_B$ among these galaxies suggests that the stellar X-ray luminosity does not scale with optical luminosity, at least for galaxies of low optical luminosities. This could be the result of a decrease in the proportion of LMXBs with decreasing optical luminosity, and/or the effects of fluctuations in the small number of LMXBs expected. ",
        "introduction": "\\label{sec:intro}  One surprising discovery uncovered by the {\\it Einstein} Observatory was that elliptical and S0 galaxies are relatively bright, extended X-ray sources (e.g., Forman et al.\\ 1979; Forman, Jones, \\& Tucker 1985; Trinchieri \\& Fabbiano 1985). It soon was established that the source of these X-rays was thermal emission from hot ($\\sim$1 keV), diffuse interstellar gas, at least for the brighter X-ray early-type galaxies. However, several unsolved problems still remain concerning the X-ray emission from elliptical and S0 galaxies. Although there is a strong correlation between the X-ray and blue luminosities in these systems, the X-ray luminosity can vary by as much as a factor of $50-100$ among galaxies with the same blue luminosity (Canizares, Fabbiano, \\& Trinchieri 1987; Fabbiano, Kim, \\& Trinchieri 1992).  Ellipticals and S0s with a low X-ray--to--blue luminosity ratio ($L_X/L_B$) also have different spectral characteristics than their X-ray bright (high $L_X/L_B$) counterparts. From {\\it Einstein} observations, Kim, Fabbiano, \\& Trinchieri (1992) found that the galaxies with the lowest $L_X/L_B$ exhibited an excess of very soft X-ray emission. {\\it ROSAT} Position Sensitive Proportional Counter (PSPC) observations of several X-ray faint galaxies have confirmed that the X-ray emission in such galaxies is not characterized by emission from hot, metal-enriched gas with a temperature of $\\sim$1 keV; a two component model, in which one component has a temperature around 0.2 keV and the other a temperature greater than a few keV provides a better fit (Fabbiano, Kim, \\& Trinchieri 1994; Pellegrini 1994). This result was confirmed by an {\\it ASCA} observation of NGC~4382 (Kim et al.\\ 1996). The source of the soft $\\sim$0.2 keV emission has been postulated to be RS CVn binary systems, M star coronae, supersoft sources, or a low temperature interstellar medium (ISM), although none of these mechanisms seems to be entirely responsible for the soft X-ray emission (Pellegrini \\& Fabbiano 1994). The hard component is thought to be emission from accreting compact X-ray binaries; this is the dominant X-ray component in most spiral galaxies (Trinchieri \\& Fabbiano 1985).  It is difficult to do detailed spectral modeling of these low $L_X/L_B$ systems because of their low X-ray count rates, even for the closest and brightest examples. Instead, in this paper we will analyze the broad band X-ray ``colors\" of elliptical and S0 galaxies. These ``colors'' are the ratios of the counts in three X-ray bands. Whereas conventional spectral fitting requires a minimum of about 1000 counts to constrain reasonably the parameters of the model, only about 100 counts are needed to determine the X-ray colors. This allows much fainter X-ray galaxies to be analyzed and compared to brighter galaxies, leading to a more comprehensive investigation into the X-ray properties of early-type galaxies. A similar technique of determining colors was used to study a large sample of X-ray galaxies with {\\it Einstein} data by Kim et al.\\ (1992). In this paper, we will determine the X-ray colors of early-type galaxies observed with the {\\it ROSAT} PSPC. The increased sensitivity of {\\it ROSAT} over {\\it Einstein} allows a larger sample to be analyzed and the X-ray colors to be determined more accurately. In addition to integrated X-ray colors, we also derive the X-ray colors as a function of galactocentric radius to explore the variation of X-ray emission with position. If the emission mechanisms are different between X-ray faint and X-ray bright elliptical and S0 galaxies, the integrated colors and color profiles should shed some light on this difference.  This paper is organized as follows. In \\S~\\ref{sec:galsamp}, we define our galaxy sample and give the relevant properties of the galaxies, and in \\S~\\ref{sec:datared}, we discuss the data analysis. The X-ray luminosities and X-ray--to--optical luminosity ratios in hard and soft bands are discussed in \\S~\\ref{sec:lxlb}. The integrated X-ray colors of the galaxies are presented in \\S~\\ref{sec:colors}, and are compared to simple models. The radial profiles of the X-ray colors are given in \\S~\\ref{sec:radcolor}. In \\S~\\ref{sec:origin}, we discuss a number of different mechanisms for the X-ray emission of elliptical galaxies of varying X-ray--to--optical luminosities. We show that the very soft component in X-ray faint galaxies is unlikely to be due to stellar sources other than LMXBs or a warm ISM in \\S~\\ref{sec:soft_alternatives}. Finally, our conclusions are summarized in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have analyzed a large elliptical and S0 galaxy sample observed with the {\\it ROSAT} PSPC. We derived X-ray luminosities and X-ray--to--optical luminosity ratios in soft and hard {\\it ROSAT} bands. We defined two X-ray colors; integrated values and radial profiles in these two colors were determined for the galaxies in our sample. We find a very large range in values (a factor of $\\sim$500) for the X-ray--to--optical luminosity ratio for galaxies in our sample, as has generally been found in the past. The X-ray colors, the color profiles, and the ratio of hard to soft band luminosities all vary as a function of $L_X / L_B$, which suggests that different emission processes are important in different galaxies. We divide the galaxies into four groups based on their X-ray--to--optical luminosity ratio: Group 1, $\\log(L^h_X/L_B) < 29.7$; Group 2, $29.7 \\le \\log(L^h_X/L_B) < 30.0$; Group 3, $30.0 \\le \\log(L^h_X/L_B) < 30.4$; Group 4, $\\log(L^h_X/L_B) \\ge 30.4$.  The X-ray brightest galaxies (Group 4) have colors and luminosities which are consistent with thermal emission from hot gas at temperatures around 0.8 keV and abundances of approximately one-half solar. These galaxies have hard and soft band X-ray luminosities which increase in proportion to one another, and which increase rapidly with the optical luminosity. They often sit at the centers of groups, and there is evidence that this increases their X-ray luminosities. The galaxies in deeper potential wells may have higher X-ray luminosities per unit blue luminosity because they have lost less gas during their early evolution, or have accreted group gas.  The intermediate X-ray luminosity galaxies (Groups 2 and 3) have colors which are generally consistent with thermal emission from lower temperature and abundance gas than the Group 4 galaxies. For these galaxies, the soft band X-ray luminosity--to--optical luminosity ratios ($L_X^s/L_B$) are nearly constant from galaxy to galaxy, while the hard band luminosity drops more rapidly than the optical luminosity. One explanation for this trend would be if the hard band luminosity were dominated by thermal emission by gas, and the soft band luminosity were dominated by stellar sources. This would require that the thermal X-ray luminosity decline rapidly with decreasing optical luminosity, as is seen for the brighter Group 1 galaxies. However, this would not explain why the soft band X-ray--to--optical luminosity ratio $L_X^s/L_B$ continues to decrease for $\\log(L_X^h/L_B) \\la 29.7$. Also, the {\\it ROSAT} and {\\it ASCA} X-ray spectra of these galaxies suggest that the bulk of the emission is thermal emission from hot gas with subsolar abundances (e.g., Davis \\& White 1996; Matsumoto et al.\\ 1997). Alternatively, the bulk of the X-ray luminosity might be thermal in origin, with $L_X^h/L_B$ decreasing as the gaseous iron abundance decreases. We show that lowering the iron abundance mainly reduces the {\\it ROSAT} hard band emission, because this emission is dominated by the iron $L$-line complex at $\\sim$1 keV.  The X-ray emission from most of the fainter Group 1 galaxies can best be described as the integrated emission from low mass X-ray binaries as was seen to be the case for the bulge of M31. In these galaxies, the ISM could have been removed by supernovae-driven winds. The X-ray colors and X-ray--to--optical luminosity ratios for all but a small subset of the Group 1 galaxies were consistent with the colors of the bulges of M31 and NGC~1291. The soft component previously seen in these galaxies was shown to be the result of the more complex spectrum of LMXBs; they are not well-represented by a simple, single temperature 5 keV bremsstrahlung model, which had been previously assumed for LMXBs. The X-ray colors of the nearby LMXB Her X-1 confirm the existence of a soft X-ray component in at least one Galactic LMXB. Individual LMXBs in the faintest galaxies should be resolvable with {\\it AXAF} for galaxies as far away as Virgo.  As supporting evidence that the soft component is from LMXBs, we show that the three most likely alternative stellar sources of the soft emission are unlikely to be important contributors. We find that M stars can match the X-ray colors of Group 1 ellipticals, but that they are probably too faint. We set a limit on the cumulative X-ray luminosity of old M stars from upper limits on the diffuse X-ray emission from globular clusters. This gives X-ray luminosities which are a factor of at least 300 lower than the soft band X-ray luminosities of Group 1 ellipticals. (However, this limit might not apply if the X-ray emission from M stars increased strongly with metallicity.) RS CVn stars have C32 colors which are too hard to match the observed colors of ellipticals, and adding hard emission from LMXBs only increases the discrepancy. Supersoft sources are too soft to explain the observed colors of ellipticals.  The brightest of the Group 1 galaxies might contain significant amounts of interstellar matter which has not been completely removed from the galaxy. These galaxies have higher C21 values and lower C32 values than the other Group 1 galaxies. This would account for their higher $L_X/L_B$ values as compared to the faintest Group 1 galaxies and the bulges. If the ISM has a temperature of $0.3-0.6$ keV and low metallicities, its emission would produce the observed colors when mixed with the right proportion of stellar emission.  Finally, we find that the faintest X-ray galaxies have hard band {\\it ROSAT} X-ray luminosities which are significantly less than expected from the integrated emission of LMXBs, as derived from brighter elliptical galaxies (Matsumoto et al.\\ 1997) and spiral galaxy bulges. This suggests that the fractional population or luminosity of LMXBs may be lower in less luminous galaxies. One way to produce such a variation would be if the slope of the initial mass function in early-type galaxies became flatter as the optical luminosity increased. Small number statistical variations in the number of LMXBs in these galaxies may also play a role."
    },
    "9804/astro-ph9804204_arXiv.txt": {
        "abstract": "We calculate the relativistic thermal bremsstrahlung Gaunt factor for the high-temperature plasma which exists in clusters of galaxies.  We calculate the Gaunt factor by employing the Bethe-Heitler cross section corrected by the Elwert factor.  The calculations in this paper are made for the fully ionized plasma for the following cases: $Z$ = 10 (Ne), 12 (Mg), 14 (Si), 16 (S), 26 (Fe).  We also calculate the Gaunt factor by using the Coulomb-distorted wave functions for nonrelativistic electrons following the method of Karzas and Latter.  By comparing the Gaunt factors calculated by these two different methods, we carefully assess the accuracy of the calculation.  We present the numerical results in the form of tables. ",
        "introduction": "High-temperature plasmas exist in the clusters of galaxies (Arnaud et al. 1994; Markevitch et al. 1994; Markevitch et al. 1996; Holzapfel et al. 1997).  Some clusters have extremely high-temperature electrons,  $k_{B} T_{e}$ = $10 \\sim 15$keV.  Relativistic expressions for the thermal bremsstrahlung emissivity have been discussed by many authors (Gould 1980; Rephaeli \\& Yankovitch 1997).  However, the relativistic expressions have been so far derived by power-series expansions.  Very recently the present authors (Nozawa, Itoh, \\& Kohyama 1998) have calculated the relativistic thermal bremsstrahlung Gaunt factor for the intracluster plasma by using the Bethe--Heitler (1934) cross section corrected by the Elwert (1939) factor for the following cases: $Z$ = 1 (H), 2 (He), 6 (C), 7 (N), 8 (O).  They have also calculated the bremsstrahlung Gaunt factor by using the Coulomb--distorted wave functions for nonrelativistic electrons following the method of Karzas \\& Latter (1961).  They have thereby assessed the accuracy of the calculations.  Their method of the calculation closely followed the work on the calculation of the inverse thermal bremsstrahlung by two of the present authors and their collaborators (Itoh, Nakagawa, \\& Kohyama 1985; Nakagawa, Kohyama, \\& Itoh 1987; Itoh, Kojo, \\& Nakagawa 1990; Itoh et al. 1991; Itoh et al. 1997).  In the present paper we will extend the calculation to heavier elements: $Z$ = 10 (Ne), 12 (Mg), 14 (Si), 16 (S), 26 (Fe).  For the sake of completeness, we will restate the formalism in this paper.  The present paper is organized as follows.  We will give formulations for the calculation of the relativistic thermal bremsstrahlung Gaunt factor in $\\S$ 2.  The numerical results will be presented in $\\S$ 3.  We will discuss the results and give concluding remarks in $\\S$ 4. ",
        "conclusions": ""
    },
    "9804/astro-ph9804174_arXiv.txt": {
        "abstract": "A simple function for the spectral power $P(\\epsilon,t) \\equiv \\nu L(\\nu) $ is proposed to model, with 9 parameters, the spectral and temporal evolution of the observed nonthermal synchrotron power flux from GRBs in the blast wave model. Here $\\epsilon = h\\nu/$m$_e$c$^2$ is the observed dimensionless photon energy and $t$ is the observing time.  Assumptions and an issue of lack of self-consistency are spelled out.  The spectra are found to be most sensitive to the baryon loading, expressed in terms of the initial bulk Lorentz factor $\\Gamma_0$, and an equipartition term $q$ which is assumed to be constant in time and independent of $\\Gamma_0$.  Expressions are given for the peak spectral power $P_p(t) = P(\\epsilon_p,t)$ at the photon energy  $\\epsilon  = \\epsilon_p(t)$ of the spectral power peak. A general rule is that the total fireball particle kinetic energy $E_0 \\sim \\Pi_0 t_d$, where $t_d \\propto \\Gamma_0^{-8/3}$ is the deceleration time scale and $\\Pi_0 \\equiv P(\\epsilon_p,t_d) \\propto \\Gamma_0^{8/3}$ is the maximum measured bolometric power output in radiation, during which it is carried primarily by photons with energy ${\\cal E}_0 = \\epsilon_p(t_d) \\propto q\\Gamma_0^4$.  This rule governs the general behavior of fireballs with different baryon loading.  Clean fireballs with small baryon loading ($\\Gamma_0\\gg 300$) are intense, subsecond, medium-to-high  energy $\\gamma$-ray events, and are difficult to detect because of deadtime and sensitivity limitations of previous $\\gamma$-ray detectors such as EGRET on {\\it CGRO}. Dirty fireballs with large baryon loading ($\\Gamma_0\\ll 300$) produce transient emissions which are longer lasting and most luminous at X-ray energies and below,  but these events are lost behind the glow of the X-ray and lower-energy background radiations except for rare serendipitous detections by pointed instruments.  The correlation between hardness and duration of loaded GRB fireballs ($100\\lesssim\\Gamma_0\\lesssim 1000$) follows from this rule. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804081_arXiv.txt": {
        "abstract": "The QSO 1345+584 has been spatially resolved by direct images and in spectral images, and has extended flux asymmetrically to the W, where its inner radio structure is seen. The brightest knots in the resolved flux correspond closely with knots in the curved radio jet, and the brightest knot has velocity of approach of some 3000 km/s with respect to the nucleus. Other parts of the line-emitting material appear to follow a systematic velocity field with values up to 1000 km/s with respect to the nucleus. The signal from the resolved continuum is not detected spectroscopically but accounts for 2/3 of the (rest UV) flux, so that it is likely to originate in hot stars. The QSO lies in or behind a compact group of galaxies of comparable brightness and irregular and knotty morphology, which probably form a dense physical group with very young stellar populations. ",
        "introduction": "The Space Telescope Imaging Spectrograph (STIS) instrument on the Hubble Space Telescope  (HST) offers new opportunities for QSO host galaxy studies. The instrument may be used to obtain direct images with the CCD which are deeper and less noisy than WFCP2, and UV images with the MAMA detectors. Use of a long slit with STIS enables spatially resolved spectra of the host galaxy, both through and away from the nucleus, with a pixel spacing of $\\sim$ 0.05\" in the visible and 0.025\" in the UV.  This paper presents the results of CCD imaging and spectroscopy of the QSO 1345+584, which has redshift 2.039. The high redshift moves rest-UV into the visible, so little observed UV flux is expected and no MAMA observations were attempted. The QSO was chosen for several reasons: it has detected extended Ly$\\alpha$ flux from ground-based observations (Heckman et al 1991), with $\\sim$1.5\" resolution; it is a radio-loud object with complex compact structure (Lonsdale et al 1993) that might be seen in the visible; and it lies in the continuous viewing zone (CVZ) of HST so that long integrations could be performed.  The QSO has a catalogued visible magnitude of 17.5 and is also known as 4C 58.27 or OP 577 (Hewitt and Burbidge 1993). Lehnert et al (1993) were unable to detect any resolved continuum flux around the QSO, in ground-based observations. Barthel et al (1988) noted that the bend angle of the structure is remarkable and suggests stong interaction with the surrounding medium. Bremer et al (1992) describe long-slit spectra in which no definite extended line emission was detected.  The observations were performed on 1997 Dec 19 as shown in Table 1. The short image exposure was taken to obtain the nucleus without detector saturation. The roll angle was determined by the CVZ window, and spectra were taken centred on the nucleus (slit A) and also offset arbitrarily to the north by 1.5\" (slit B), through a 2\" wide slit. The wide slit degrades the spectral resolution but was judged necessary to optimise detection of the continuum signal from the faint QSO host galaxy. All observations were performed with 3 CCD readouts, to identify cosmic rays. Standard wavelength calibration exposures were taken with the spectra.  All the images were processed with standard CALSTIS reduction, but using a mean dark image from the day of observation to eliminate hot pixels as well as possible, since the signal levels of interest are low. The individual reads were combined to eliminate cosmic rays in the standard way. We discuss below the results and measurements made in the processed data. ",
        "conclusions": ""
    },
    "9804/astro-ph9804048_arXiv.txt": {
        "abstract": "In the interstellar medium, as well as in the Universe, large density fluctuations are observed, that obey power-law density distributions and correlation functions. These structures are hierarchical, chaotic, turbulent, but are also self-organizing. The apparent disorder is not random noise, but can be described by a fractal, with a deterministic fractal dimension.  We discuss the theories advanced to describe these fractal structures, and in particular a new theory of the self-gravity thermodynamics, that could explain their existence, and predict their fractal dimension. The media obeying scaling laws can be considered critical, as in second order phase transitions for instance. ",
        "introduction": "Fractals have been introduced by Mandelbrot (1975) to define geometrical ensembles, or mathematical sets, that have a fractional dimension. He pioneered the study of very irregular mathematical sets, where the methods of classical calculus cannot be applied. Fractals are not smooth nor differentiable; they are characterized by self-similarity. Their geometrical structure has details at all scales, and the details are representative of the whole.  Fractals are very rich, since they give the best approximation for many natural phenomena, that cannot be represented by regular and smooth geometries (natural systems are only fractal between two boundaries, with upper and lower cut-offs, and not true mathematical fractals, of course). There can exist a large part of randomness in fractals, and the fractal dimension and exponents of scaling laws are the tools to quantify the hidden order in them.  There are many definitions of the dimension, trying to quantify how much space a system fills. We will use mainly the Hausdorff dimension $D$, based on the Hausdorff measure, that generalizes the notion of length, area and volume. When the length scale is magnified by $\\lambda$, a regular and smooth set's area and volume are respectively scaled by $\\lambda^2$, $\\lambda^3$, while the fractal set measure is scaled by $\\lambda^D$. For the astrophysical systems that we describe below, the mass contained within a scale $r$ is M$\\propto r^D$, with $D$ a fractional number between 1 and 3.   \\bigskip ",
        "conclusions": "We have emphasized the existence of two astrophysical fractals, the interstellar medium, with structures ranging from 10 AU to 100 pc, and the large-scale structures of galaxies, from 50 kpc to 150 Mpc at least. The first one is in statistical equilibrium, while the second one is still growing to larger scales. In both cases, we can describe these media as developping large-scale fluctuations with large correlations as is familiar in critical phenomena. We propose that in both cases, self-gravity is the main force governing these fractal structures.  The statistical thermodynamic approach is developped, and it is shown that the phenomenological potential, which is in $1/r$ between two cutoffs (at small and large-scale), can be described by a scalar field theory. We use the renormalization group methods for this scale-independent system, to find the universality class of the problem. The fractal dimension, and the potential correlations exponent can be derived from the critical exponents of the Ising $d=3$ model. The stability of the results with respect to perturbations has been studied; the results are quite robust with respect to perturbations of external forces. Of course,  if the external energy is large with respect to the self-gravitational energy of the gas (in the vicinity of violent star-formation in the ISM for instance), then the hierachical structure will be destroyed. The observed fractal dimensions are compatible with the predictions, and other observational tests are proposed.   \\bigskip"
    },
    "9804/astro-ph9804013_arXiv.txt": {
        "abstract": " ",
        "introduction": "The luminosity function (LF) $\\phi (L)$ of galaxies is defined as the number density of galaxies per unit luminosity $L$: the number of galaxies in volume ${\\rm d}V$ with luminosity between $L$ and $L + {\\rm d}L$ is $\\phi(L) {\\rm d}L {\\rm d}V$.  The LF represents one of the most direct interfaces between studies of dwarf galaxies and cosmology, the two subjects of this conference. This statement follows directly from the form of the fluctuation spectrum -- most plausible theories of structure formation (e.g.~cold dark matter [3]) predict a fluctuation spectrum whose power spectrum $\\vert \\delta (k) \\vert^{2} \\propto k^n$ has index $n \\sim -2$ on scales of galaxies.  This value of $n$ simultaneously requires that low-mass systems be far more numerous than high-mass ones, so that the statistical properties of low-mass systems are a powerful probe of galaxy formation theories.  Consequently one of the most fundamental predictions of models of galaxy formation and evolution is the LF, particularly at faint magnitudes (e.g.~refs.~6,13,48--50). These models generally assume a power spectrum of primordial dark matter fluctuations (say, from CDM theory) and then adopt prescriptions for the behavior of the baryonic component (specifically, things like the effects of gas dissipation, star-formation efficiencies, feedback from supernovae, and the temperature-density structure of the intergalactic medium need to be quantified).  Significant differences in the theoretical luminosity functions from different models are seen, even when similar dark matter power spectra are assumed, because of differing prescriptions for the behavior of the baryonic component.  For example Babul \\& Ferguson [2] suggest a faint-end slope of $\\alpha \\sim -2.6$ at faint magnitudes, whereas White \\& Kauffmann [49] compute values more like $\\alpha \\sim -1$ for their dwarf suppression models (here $\\alpha$ is the logarithmic slope of the LF: $\\alpha = {\\rm d} \\log \\phi(L) / {\\rm d} \\log L$).  Therefore measuring $\\alpha$ as accurately as possible at faint magnitudes seems like a worthwhile exercise.   Clusters of galaxies have proved a popular environment for determining the LF.  This is because the distances to the galaxies there are known (at least in a statistical sense) so that the LF can be determined by photometry alone and spectroscopic redshifts are not needed.  This in turn means that the LF can be measured down to very faint limits.  But there is a price to be paid for the cosmologist.  Cluster galaxies are in dense environments where the crossing time plus age is less than or comparable to the Hubble time.  These galaxies have therefore had their properties shaped by cluster-related processes e.g.~ram-pressure stripping of gas from the galaxies.  Measurements of the galaxy LF in clusters might then tell us more about these cluster-related processes than about cosmological parameters like the shape of the power spectrum on small mass scales. This is particularly true of very dense environments like the centers of rich clusters with large fractions of elliptical galaxies.  It is less true of diffuse spiral-rich environments and the outer parts of clusters.  In Figure 1, for a sample of clusters with LF determinations, I present the velociy dispersion (an indication of the mass of the cluster) as a function of the cluster redshift $z$.  Also shown on this figure are lines corresponding to 200 kpc = 10 arcmin -- this is where a rich cluster core roughly fits into a single 2K CCD frame on a 2 to 4 meter telescope -- and corresponding to where a typical dSph galaxy with $R=24$ has a scale-length (as inferred from the Binggeli plot) of 1 arcsec -- to the right of this line, most faint cluster galaxies detected have intrinsic scale-lengths smaller than the seeing and so cannot be readily discriminated from background galaxies on morphological grounds.  \\begin{figure} \\vspace{300pt} \\caption{ The velocity dispersions and redshifts for some clusters with recent LF determinations} \\end{figure}  The techniques used to determine the LF in a particular cluster are determined by its position on Figure 1.  In simplest terms, clusters are either on the left of the dashed lines (hereafter I call these ``nearby clusters'') or between or to the right of the dashed lines (I call these ``distant clusters'').  Most early work in this subject that reached absolute magnitudes faint enough to probe dwarf galaxies (those in sequence (4) of the Binggeli diagram) were photographic measurements of nearby clusters e.g.~Virgo [29] and Fornax [10]. Clusters significantly more distant than Coma (see e.g.~ref.~35 for a large photographic survey of Coma) could not be studied in this way because (i) photographic plates were not sensitive enough to probe the bulk of the dwarf population, and (ii) cluster members could not be readily distinguished from background/foreground galaxies.  However, the advent of 2K CCDs on 2 -- 4 m telescopes allowed very precise statistical determinations of field galaxy counts down to very faint magnitudes (e.g.~refs.~43 and 17). Background/foreground contamination in cluster fields could then be estimated and corrected for down to $R=25$, and the LFs of several rich clusters out to $z = 0.2$ were measured. The main limitation of these studies was that the errors at the faintest magnitudes were quite large because of uncertainty in the field-to-field variance of the background.  However, large samples of clusters could be studied, and these errors reduced when combining samples of similar clusters.  In Section 5, some results of this type will be presented and discussed.  These 2K CCDs are not well suited to studies of nearby clusters, although individual dwarfs can easily be identified and no background subtraction is needed.  This is because the CCDs are small and each exposure can only cover a negligible fraction of the cluster core (we are to the left of the 200 kpc = 10 arcmin line in Figure 1).  Counting statistics are then crippling -- a typical 2K CCD field in the center of the Virgo Cluster will have no or at most 1 or 2 dwarfs brighter than $R=22$.  More recently, 8K and larger CCDs have become available on 4 m telescopes (e.g.~the UH 8K mosaic on the 3.6 m CFHT [23]). Large samples of galaxies in nearby clusters can now be obtained and the LFs measured accurately down to very faint magnitudes. Because no background subtraction is required, the only errors at faint magnitudes come from counting statistics -- the errors are consequently much smaller than the errors at the faint end for the distant clusters, at least for reasonably dense nearby clusters like Virgo and Fornax. An alternative approach has been employed by Phillipps et al.~[27] who use sensitive median-stacked UK Schmidt plates [32] to measure the Virgo cluster LF down to $M_R \\sim -11.5$ (see the paper by Jones et al.~in these proceedings -- this was one of the most significant new results at this meeting). A project in collaboration with B.~Tully (University of Hawaii) has been undertaken to measure the Ursa Major LF but this cluster is sufficiently diffuse that counting statistics are still a significant problem. ",
        "conclusions": ""
    },
    "9804/astro-ph9804225_arXiv.txt": {
        "abstract": "We report the discovery of two simultaneous quasi-periodic oscillations in the persistent flux of GX 349+2 at frequencies $712\\pm 9$ and $978 \\pm 9$ Hz, with $rms$ amplitudes 1.25\\%$\\pm$0.34\\% and 1.34$\\pm$0.32\\%, respectively.  During our 152 ks observation with the Rossi X-ray Timing Explorer, GX 342+2 was in either the normal branch or the flaring branch with count rates in the nominal 2-60 keV RXTE-PCA band ranging from a low of 8,000 cps to a high of 15,000 cps.  The kHz QPOs were observed only when the source was at the top of the normal branch when the count rate was about 8,200 cps corresponding to a flux of $1.4\\times 10^{-8}$ ergs cm$^{-1}$ s$^{-1}$ in the 2-10 keV band. With this report, now kHz QPOs have been observed in all the 6 Z sources. ",
        "introduction": "GX 349+2, also known as Sco X-2, is one of 6 bright low mass X-ray binaries whose X-ray color-color diagrams resemble the shape of the letter Z, which, therefore, have been referred to as the Z sources (\\cite{hk89}). They are among the brightest X-ray sources in the sky. Their time variability in the frequency domain below 100 Hz as charatcerized by their FFT power spectra is closely correlated with their energy spectral state as characterized by their X-ray color-color diagram. Since the launch of the Rossi X-Ray Timing Explorer (RXTE), quasi-periodic intensity oscillations with frequencies in the range of several hundred hertz to over one thousand hertz have been discovered in 5 of the 6 Z sources:  Sco X-1 (\\cite{scox1_discovery}), Cyg X-2 (\\cite{cygx2_discovery}), GX 17+2 (\\cite{gx17+2_discovery}), GX 5-1 (\\cite{gx5-1_discovery}), and GX 340+0 (\\cite{cygx2_discovery}). GX 349+2 is the only source from which no kHz QPOs have been reported to date.  In this letter we report the discovery of kHz QPOs in GX 349+2 and compare their characteristics with those of the QPOs observed in the other Z and atoll sources. ",
        "conclusions": "We have reported the discovery of two simultaneous QPO peaks in the FFT power spectrum of GX 349+2. In this section we place our results in the context of the kHz QPO phenomenologies of the other Z sources.  GX 349+2 is the last Z source from which such characteristic twin QPO peaks have been found. The color-color diagram we observed is also quite characteristic of those observed in the past.  Compared to other Z sources, GX 349+2 has never been observed to be on its horizontal branch.  Four of the 6 Z sources, i.e.,  Cyg X-2, GX 5-1, GX 17+2, and GX 340+0, have clearly observable horizontal branches. The kHz QPO characteristics of those sources are nearly identical in that they all show up on the horizontal branch of the Z diagram. As the source moves along the Z track toward the vertex of the horizontal branch and the normal branch and further down the normal branch, the QPO centroid frequencies increase and their $rms$  amplitudes decreases until they cease to be observable in the middle part of the normal branch.  The frequency difference of the twin QPO peaks are consistent with being constant for all of these sources. The differences are, respectively, 343$\\pm$21, 327$\\pm$11, 294$\\pm$8, and 374$\\pm$24 Hz.  In the case of Sco X-1, the kHz QPO peaks are observed on the normal branch.  The separation of the two QPO peaks changes from a high of 292$\\pm$2 to a low of 247$\\pm$3 Hz as Sco X-1 moves down the normal branch toward the flaring branch (\\cite{scox1_separation}).  The horizontal branch, if any, of the Z track of the GX 349+2 has not been observed.  The kilohertz QPO characteristic we have observed in GX 349+2 is very similar to those of Sco X-1 and the other Z sources in that the kHz QPOs exist on the upper or middle normal branch. As the source approaches the flaring branch, the kilohertz QPOs become weak and cease to be observable.  Unfortunately, our existing data from GX 349+2 are not sufficient in quantity to further investigate the characteristics of the observed QPOs. Further and more extensive observations, especially when GX 349+2 is at the upper normal branch, are highly desirable and encouraged."
    },
    "9804/astro-ph9804155_arXiv.txt": {
        "abstract": "An analysis of some properties of stellar models as a function of metallicity $Z$ (and helium content $Y$) is presented, with special attention to those stars with metallicities higher than twice or three times solar.  It is shown that the stellar properties as a function of $Z$ are mainly determined by the effects of the opacities at sub-solar metallicities, and by the effects of the mean molecular weight and stellar mass loss at higher metallicities. As a result, very metal-rich stars ($Z\\simgr 0.05$) exhibit properties that deviate from what is expected from the known characteristics at lower metallicities. In particular, they are more luminous and hotter than those at $Z\\simlr 0.05$ due to the effect of the mean molecular weight. They have main sequence lifetimes much shorter (60\\% shorter at $Z=0.1$ than at $Z=0.02$) than those at solar metallicity due to their lower initial hydrogen content. Finally, the high mass loss rates at high metallicities affect significantly the population synthesis of massive stars in very metal-rich regions. An analysis of expected Wolf-Rayet and supernovae populations in such conditions is briefly presented. ",
        "introduction": "Observations as well as models of galactic chemical evolution suggest that the metallicity at the center of our galaxy, where the density of matter is high, can reach up to 3 to 5 times the solar metallicity (McWilliam \\& Rich 1994, Simpson et al. 1995). Other objects such as elliptical galaxies, or maybe even quasars (Korista et al. 1996), also show evidences of very metal-rich stellar populations.  The metallicity affects the evolution of stars mainly through its impact on the radiative opacities, the equation of state, the nuclear reaction rates and the stellar mass loss rates. The impact of the metallicity on the stellar properties has been well studied for $0.001\\le Z \\le 0.04$ on base of extensive stellar model calculations (see e.g. Schaller et al. 1992, hereafter Paper I, and Maeder \\& Conti 1994). In contrast, they are still poorly known at metallicities much higher than solar. Recent calculations have been performed at metallicities five times solar (Fagotto et al. 1994, Mowlavi et al.  1997). Mowlavi et al. (1997) notice that those very metal-rich models present specific properties which are distinct from the ones at $Z\\le 0.04$. In particular, they are {\\it hotter} and {\\it more luminous} that those at solar metallicity, contrary to expectations. It is the purpose of this paper to analyze in detail some of these properties of very metal-rich stars. These could have important consequences on the observable properties of high metal-rich galaxies.  The study is based on the extensive grids of stellar models computed by the Geneva group with $0.001 \\le Z \\le 0.1$. It covers the main phases of stellar evolution, excluding the latest stages of evolution. The core helium burning phase of very metal-rich low-mass ($M<\\mmass{1.7}$) stars, important for the study of elliptical galaxies, is not considered in this paper as it has been the subject of existing papers in the literature (see, for example, Dorman, Rood \\& O'Connell, 1993, Fagotto et al. 1994).  A general discussion on the dependence of stellar properties on chemical composition is presented in Sect.~\\ref{Sect:chemical composition}. Section \\ref{Sect:models} summarizes the model ingredients used in the computation of the Geneva grids. The stellar properties during the main sequence (MS) and the core helium burning (CHeB) phase are analyzed in Sects.~\\ref{Sect:MS} and \\ref{Sect:CHeB}, respectively. Wolf-Rayet stars in very metal-rich environments are briefly discussed in Sect.~\\ref{Sect:WR}.  Conclusions are drawn in Sect.~\\ref{Sect:conclusions}. ",
        "conclusions": "\\label{Sect:conclusions}  This paper analyzes some properties of stellar models as a function of the metallicity, with particular attention to those at very high metallicities (above twice or three times solar). The stellar properties as a function of $Z$ are shown to result from the action of four ingredients\\footnote{We neglect in this paper the effects of mild mixing processes due for example to rotation, diffusion, gravitational settling, etc... which may modify the internal composition profiles and some of the values given here}: the mean molecular weight, an increase of which leads to more luminous and hotter stars (the $\\mu$-effect); the opacity, an increase of which leads to less luminous and cooler stars (the $\\kappa$-effect); the nuclear energy production (but to a lesser extent), an increase of which leads to less luminous and cooler stars with lower $T_c$ and $\\rho_c$ (the $\\varepsilon_{nuc}$-effect; and mass loss, which acts indirectly on the stellar structure by reducing the stellar mass (dramatically in the most massive metal-rich stars). The $\\kappa$-effect plays the main role at $Z\\simlr 0.04$, while the $\\mu$-effect becomes preponderant at $Z\\simgr 0.02$. In the later case, the $\\Delta Y/\\Delta Z$ law is important.  Based on these considerations and on the Geneva's grids of stellar models, it is shown in particular that very metal-rich stars exhibit properties which cannot be extrapolated from model characteristics at lower metallicities. In particular,  \\noindent - very metal-rich stars (such as at $Z=0.1$) are more luminous and hotter than those at $Z\\le 0.05$. This property contrasts with the usually known trend at lower metallicities (whereby $L$ and $T_{eff}$ decrease with increasing $Z$ due to the $\\kappa$-effect);   \\noindent - very metal-rich stars have much lower MS lifetimes than those at $Z\\le 0.05$. For example the MS lifetime at $Z=0.1$ is 2.5 times shorter than that at solar metallicity, independent of their initial mass (at least up to \\mass{40});  \\noindent - stellar population synthesis in metal-rich environments is significantly affected by the increased mass loss rates at high $Z$. In particular, a) massive stars with $Z=0.1$ enter the WNL Wolf-Rayet phase very early during their evolution, and those more massive than about \\mass{50} probably never enter the WC phase; b) the most massive ($M\\ge\\mmass{60}$ at $Z=0.1$) high-metallicity stars loose the majority of their mass during their MS phase, and probably end their life as He white dwarfs; and c) the number of WR stars over O-type stars increases with metallicity, and metal-rich regions are expected to be characterized by a high fraction of type I (b/c/other?) supernovae relative to the total number of supernovae having massive star progenitors (up to 1 out of 3 core collapse supernovae events at $Z=0.1$)."
    },
    "9804/astro-ph9804319_arXiv.txt": {
        "abstract": "We examine the formation of giant elliptical galaxies using their globular cluster (GC) systems as probes of their evolutionary history. The bimodal distributions of GC metallicities in such galaxies are often cited as evidence for the formation of giant elliptical galaxies through mergers involving gas-rich spirals, with the metal-rich GCs forming during the merger process. We explore an alternative possibility: that these metal-rich clusters represent the galaxy's intrinsic GC population and that the metal-poor component of the observed GC metallicity distribution arises from the capture of GCs from other galaxies, either through mergers or through tidal stripping. Starting with plausible assumptions for the luminosity function of galaxies in the host cluster and for the dependence of GC metallicity on parent galaxy luminosity, we show using Monte-Carlo simulations that the growth of a pre-existing seed galaxy through mergers and tidal stripping is accompanied naturally by the capture of metal-poor GCs whose chemical abundances are similar to those which are observed to surround giant ellipticals.  We also investigate the spatial distribution of GCs in isolated galaxies of low and intermediate luminosity and conclude that, at the epoch of formation, the GCSs of such galaxies are likely to have been more spatially extended than their constituent stars. Thus, the capture of GCs through tidal stripping, unlike mergers, does not necessarily conserve GC specific frequency. Comparisons of model GC metallicity distributions and specific frequencies to those observed for the well-studied galaxies NGC 4472 (=M49) and NGC 4486 (=M87), the two brightest cluster members of the nearby Virgo cluster, show that {\\it it is possible to explain their bimodal GC metallicity distributions and discordant specific frequencies without resorting to the formation of new GCs in mergers or by invoking multiple bursts of GC formation.} Finally, we discuss the possibility of using the ratio of the number of metal-poor and metal-rich GCs in giant elliptical galaxies as a diagnostic of their merger histories. We use this method to derive upper limits on the number of galaxies and total luminosity accreted to date by NGC 4472. ",
        "introduction": "\\bigskip  An understanding of galaxy formation is one of the principal goals of modern astrophysics. An important first step in this campaign came with the demonstration by Toomre (1977) that mergers of late-type galaxies often produce end-products which bear a striking resemblance to giant ellipticals (gEs). In subsequent years, support has grown for the notion that mergers play an important, if not dominant, role in the formation of gE galaxies. Not only have investigations into the structure and kinematics of low-redshift gEs revealed fossil evidence of past mergers (Bender \\& Surma 1992; Tremblay \\& Merritt 1996; Faber et al. 1997), but theoretical models of galaxy formation invariably predict the growth of such galaxies through mergers (Kauffmann, White \\& Guiderdoni 1993; Cole et al. 1994; Navarro, Frenk \\& White 1995). Particularly close attention has been paid to cD galaxies and centrally dominant ellipticals in rich clusters since it is widely believed that the formation of these rare objects is fundamentally connected to the formation of the clusters themselves ($e.g.$, Ostriker \\& Tremaine 1975; Hausman \\& Ostriker 1978; Merritt 1984, 1985; Mackie 1992; West 1994).  The globular cluster systems (GCSs) of gE galaxies have traditionally provided some of the most stringent constraints on models for the formation and evolution of these objects ($e.g.$, Harris 1986). Since GCs are the oldest and brightest objects which can be identified in gE galaxies, they are powerful tools for investigating the structure and chemical evolution of their host galaxies at the epoch of formation. Indeed, any model which seeks to explain the formation of gEs must be able to account for the {\\it number}, {\\it chemical abundance} and {\\it spatial distribution} of their associated systems of GCs.  One of the strongest arguments against the view that gE galaxies represent the remains of merged spirals ($e.g.$, Toomre 1977) is that of van den Bergh (1982). Simply stated, ellipticals have GC specific frequencies ($i.e.$, the number of GCs per unit parent galaxy luminosity; see \\S2.3) which are two to three times higher than those of spirals and irregulars (Harris 1991). Of course, these morphologically disparate galaxies exhibit very different clustering properties, so the extent to which their GCSs have been influenced by environment is unclear (see West 1993). For instance, Forte, Martinez \\& Muzzio (1982) suggested that the high specific frequency of NGC 4486, the centrally dominant galaxy in the Virgo cluster, might be explained by the capture of GCs through tidal-stripping of less massive galaxies (c.f. Muzzio 1987).  An additional constraint on galaxy formation models comes from the observed distribution of GC metallicities. Following initial suggestions of non-Gaussian metallicity distributions based on rather small samples of GCs (Ostrov, Geisler \\& Forte 1993; Lee \\& Geisler 1993), a flurry of activity has now demonstrated beyond question that many, and perhaps {\\it most}, luminous ellipticals show distinctly bimodal GC metallicity distributions ($e.g.$, Whitmore et al. 1995; Elson \\& Santiago 1996). The number of gEs having reliable GC metallicity distributions has now risen to roughly a dozen (see the recent compilation of Forbes, Brodie \\& Grillmair 1997; hereafter FBG97). As these authors point out, the mean metallicity of the {\\it metal-rich} component increases with the luminosity of the parent galaxy in a manner which is entirely consistent with that originally proposed by van den Bergh (1975). However, the position of the {\\it metal-poor} peak shows no strong correlation with the properties of the host galaxy: it falls in the range $-1.5 \\lae {\\rm [Fe/H]} \\lae -0.7$ with little or no dependence on parent galaxy luminosity. The challenge for models of gE formation is therefore to explain not just the origin and size of these chemically-distinct GC populations, but also to understand why the location of the metal-poor component peak depends so weakly on the luminosity of the parent galaxy.  One possible explanation for the high specific frequencies and bimodal GC metallicity distributions of gEs is that the same mergers which give rise to the {\\it galaxies themselves} also serve as catalysts for the formation of new GCs (Schweizer 1986; Ashman \\& Zepf 1992). In other words, the populations of metal-rich GCs associated with galaxies such as NGC 4486 (Whitmore et al. 1995; Elson \\& Santiago 1996) and NGC 4472 (Geisler, Lee \\& Kim 1996) are interpreted as being the relics of merger-induced GC formation, whereas the metal-poor components are assumed to constitute the {\\it pre-existing} GCSs. However, as Geisler, Lee \\& Kim (1996) point out, both of these galaxies show bimodal GC metallicity distributions yet have specific frequencies which differ by nearly a factor of three (Harris 1986), so it seems unlikely that any {\\it single} mechanism, such as merger-induced GC formation, can explain entirely the properties of their GCSs. In light of these difficulties, FBG97 have argued that the chemically-distinct GC populations in gEs are indigenous to their parent galaxies and are formed in ``two distinct phases of star formation from gas of differing metallicity\". Nevertheless, it is difficult to identify the mechanism responsible for the separate GC-formation episodes and for delaying the onset of the second burst.  In this paper, we show that bimodal GC metallicity distributions in gEs are a natural consequence of the capture of GCs through mergers and tidal-stripping of other galaxies. Our model differs from that of Ashman \\& Zepf (1992) in that we identify the {\\it metal-rich} component of the GC metallicity distributions as the gE's intrinsic GCS; mergers and tidal stripping will lead almost inevitably to a second peak in the range $-1.5 \\lae {\\rm [Fe/H]} \\lae -0.7$, with the exact position depending weakly on the shape of the luminosity function in the host cluster. We also conclude that, contrary to some earlier claims, the capture of GCs through tidal-stripping of other cluster galaxies does not necessarily conserve specific frequency, particularly if the bulk of the tidally-stripped GCs originated in low- and intermediate-luminosity galaxies. Finally, we use this model to explain the bimodal GC metallicity distributions and discordant specific frequencies of NGC 4472 and NGC 4486 --- the two brightest members of the Virgo cluster. ",
        "conclusions": "Many, perhaps most, gE galaxies have GCSs which show bimodal metallicity distributions.  We argue that such bimodal distribution may arise from the capture of other GCs, either through mergers or tidal stripping. For reasonable choices for the LF of galaxies in the host cluster and for the dependence of GC metallicity on parent galaxy luminosity, such captured GCs define a metal-poor population which has a maximum in the range $-1.5 \\lae {\\rm [Fe/H]} \\lae -0.7$. Since the precise location of this maximum depends primarily on the LF of the surrounding galaxy cluster, mergers and tidal stripping provide a natural explanation for not only the origin of the metal-poor peak, but also its poor correlation with the luminosity of the gE (FBG97). {\\it Our principal conclusion is that it is possible to explain bimodal GC metallicity distributions without resorting to the formation of new GCs in mergers or by invoking multiple bursts of GC formation.}  A comparison between the simulated and the observed GC metallicity distributions for the well-studied galaxy NGC 4472 shows excellent agreement, particularly for steeper LFs. We argue that the weaker cluster tidal field near NGC 4472 compared to NGC 4486 is responsible for the disparity in the sizes of the GCSs of these two galaxies since tidal stripping may, in some cases, lead to substantial increases in specific frequency. This is particularly true if the metal-poor GCs in gEs were captured through tidal stripping of low- and intermediate-luminosity galaxies since the GCSs of such galaxies are likely to have been more extended than their constituent stars at the epoch of formation.  It would seem the time is ripe to re-examine the dynamical evolution of GCSs using improved numerical simulations ($i.e.$, existing N-body experiments suffer from limited resolution and restricted dynamic range; Muzzio 1987). With larger simulations and more sophisticated algorithms it will be possible to investigate the time evolution of the size, spatial extent and metallicity distribution of the GCSs associated with bright cluster ellipticals.  We describe a new method for placing limits on the fraction of luminous matter which has been captured by individual gE galaxies. An application of the method to NGC 4472 suggests that as much as $\\sim$ 65\\% of its present-day luminosity could have been acquired in this way. This technique may provide a useful tool for studying the merger histories of gE galaxies."
    },
    "9804/astro-ph9804069_arXiv.txt": {
        "abstract": "We perform  numerical simulations of  galaxy mergers between star-forming and gas-rich spirals in order to explore the origin of the Fundamental Plane (FP) of elliptical galaxies. We consider particularly  that the origin of the slope of the FP  is essentially due to the non-homology in structure and kinematics of  elliptical galaxies and  accordingly investigate structural and kinematical properties of elliptical galaxies formed by dissipative galaxy merging with star formation. We found that the rapidity of star formation, which is defined as the ratio of dynamical time-scale of merger progenitor to the time-scale of gas consumption by  star formation, is a key determinant for nonhomology parameters, such as the density profile of stellar component, the relative importance of global rotation in kinematics, and the ratio of total dynamical mass to luminous mass, in merger remnants. We furthermore found that this result does not depend so strongly on initial intrinsic spins of progenitor disks and orbital energy and angular momentum of mergers. These  results strongly  suggest that the structural and kinematical nonhomology observed in elliptical galaxies can be  closely associated with the difference in  star formation history between elliptical galaxies formed by dissipative galaxy merging. Based upon these results, we discuss a  close physical relation between the origin of the FP and the star formation history of elliptical galaxies. ",
        "introduction": "The Fundamental Plane (FP) is one of the most important universal relations in early type galaxies because it contains valuable information about the formative and evolutionary process  of galaxies (\\cite{dd87,dre87}). The commonly used form of the scaling relation in the FP is described as $R_{\\rm e}$= $\\sigma^A I^B$, where $R_{\\rm e}$, $\\sigma$, and $I$ are effective radius, central velocity dispersion, and mean surface brightness of elliptical galaxies, respectively. The exponents $A$, $B$ are considered to be $1.56 \\pm 0.07$ and $-0.94 \\pm 0.09$ in the FP derived by $K$ band photometry, respectively, and these values  deviate significantly from the  values $A = 2.0$ and $B = -1.0$ expected from virial theorem (e.g., \\cite{pah95}). This apparent deviation requires that the ratio of dynamical mass ($M$) to luminosity of elliptical galaxies  ($L$)  depends on $M$ as $M/L$ $\\propto$ $M^{\\alpha}$ ($\\alpha$ = 0.12 $\\pm$ 0.03 for $K$ band). To understand the physical meaning of this dependence is one of the fundamental issues in the theory of galaxy formation (e.g, Bender, Burstein, \\& Faber 1992, 1993). Possible interpretations for the required dependence of $M/L$ on $M$ (the ``slope'' or ``tilt'' of the FP) are generally considered to be divided into the following two. One is that the required dependence of  $M/L$ on $M$ results from the fact that the mean stellar age and metalicity of elliptical galaxies depend systematically  on  $M$. The other is that the required dependence  reflects the $M$  dependence of   structural and kinematical properties of elliptical galaxies (``nonhomology''). Although we should not neglect the importance of stellar populations in generating the $M$  dependence  of the $M/L$ (Renzini \\& Ciotti 1993; Pahre \\& Djorgovski 1997), we here consider that the origin of the  required $M$   dependence of $M/L$ is closely associated with the structural and kinematical properties dependent on $M$ or $L$ in elliptical galaxies.   A growing number of observational studies have been accumulated which reveal the nonhomology in structural and kinematical properties of elliptical galaxies, such as the  systematical departure from the universal light profile of elliptical galaxies, the so-called $R^{1/4}$ law (e.g., Burkert 1993; Caon, Capaccioli, \\& D'Onofrio 1993), the difference in the relative importance of global rotation in kinematics between elliptical galaxies (e.g., Davies et al. 1983), and the ratio of dynamical mass to luminous mass (e.g., Rix 1997). These nonhomology observed in  elliptical galaxies are considered to be one of the most promising candidates which can provide  a plausible explanation for the origin of the slope of the FP (Ciotti, Lanzoni, \\& Renzini 1996; Pahre et al. 1995; Prugniel \\& Simien 1996; Graham \\& Mathew 1997). There are a number of  theoretical attempts to provide  a physical basis for the origin of the above structural and kinematical nonhomology of elliptical galaxies. Hjorth \\& Madsen (1995) constructed a self-consistent equilibrium model of elliptical galaxies based upon the statistical mechanics of violent relaxation and demonstrated that if the more luminous elliptical galaxies have the so-called  $R^{1/n}$ light profile with larger $n$, the slope of the FP can be successfully reproduced. Capelato, de Carvalho, \\& Carlberg  (1995) performed  a set of numerical simulations on  dissipationless galaxy mergers with variously different initial conditions of mergers and found that initial orbital energy of merging systems is a primarily important factor for the radial profile of velocity dispersion and mass distribution of merger remnants and thus for the successful reproduction of  the FP slope. These studies have indeed provided clues about  the understanding of the origin of the FP slope, it has not been clarified,  however,  why and how the  exponent  $n$ in the $R^{1/n}$ light profile and the initial orbital energy of merging systems depend primarily on the galactic luminosity or mass. Thus it is vital to make it clear why and how the structural and kinematical nonhomology required for the reproduction of the FP slope depends on the galactic luminosity or mass.    The purpose of this paper is to elucidate the origin of the luminosity-dependent structural and kinematical properties (``nonhomology'') in elliptical galaxies and thereby to explore the physical meaning of the FP. We here focus particularly on the star formation history of elliptical galaxies  and accordingly investigate the important roles of the star formation history in generating the structural and kinematical nonhomology of elliptical galaxies. In the present study, the origin of the  luminosity-dependent structural and kinematical properties (nonhomology) is considered as follows. First, the difference in structural and kinematical properties of elliptical galaxies is determined primarily by the difference in the star formation history of elliptical galaxies. Next, the star formation history of forming elliptical galaxies depends  principally on galactic luminosity. The structural and kinematical properties of elliptical galaxies therefore depend on galactic luminosity. This consideration for the understanding of the origin of the FP slope is reasonable and realistic, since difference in star formation history between elliptical galaxies and its possible dependence on galactic luminosity are suggested both by theoretical studies and by observational ones  on  chemical and photometric evolution of elliptical galaxies such as  color-magnitude relation (e.g., \\cite{ay87}), the abundance ratio [Mg/Fe] (\\cite{wor92,mat94,ben96}), and $H_{\\beta}$ line index (Faber et al. 1995; Worthey, Trager, \\& Faber 1996). In investigating the non-homologous nature in elliptical galaxies, we adopt the ``merger hypothesis'' in which elliptical galaxies are formed by galaxy mergers between two late-type spirals (Toomre \\& Toomre 1977). In particular, we consider galaxy mergers between two disk galaxies with a gas mass fraction larger than 0.2 observed in the present typical late-type spirals in order to mimic elliptical galaxy formation by galaxy mergers at higher redshift, which is suggested by recent observational studies on the tightness of the color-magnitude relation and the FP and on the photometric evolution of elliptical galaxies (Bower et al. 1992; Ellis et al. 1996; Franx \\& van Dokkum 1996). We do not here intend to discuss  formation of elliptical galaxies within a framework of a specific cosmology such as the CDM model (e.g., White \\& Rees 1978; White \\& Frenk 1992; Baugh, Cole, \\& Frenk 1996).     The  layout of this paper is as follows. In \\S 2, we summarize numerical models used in the present study and main points of analysis of the study. In \\S 3, we describe the dependence of  structural and kinematical properties of merger remnants on the star formation history of galaxy mergers. In this \\S 3, we furthermore discuss how the  star formation history of galaxy mergers should depend on galactic luminosity for the successful reproduction of the FP slope observed in elliptical galaxies. In \\S 4, we discuss a close relation between the origin of the slope and the tightness of the FP and the star formation history of galaxy mergers. The conclusions of the present study are given in \\S 5. ",
        "conclusions": "We have numerically investigated  structural and kinematical properties of  remnants of galaxy mergers between late-type  spirals considerably abundant in interstellar gas in order to explore the origin of the fundamental plane (FP) of elliptical galaxies. We particularly considered  that the origin of the slope of the FP  is essentially due to the nonhomology in structure and kinematics of  elliptical galaxies formed by dissipative galaxy merging with star formation. We found that the rapidity of star formation, which is defined as the ratio of dynamical time-scale of merger progenitor to the time-scale of gas consumption by  star formation, is a key determinant for nonhomology parameters such as the density profile, the degree of global rotation, and the ratio of total dynamical mass to luminous mass, in merger remnants. This result strongly suggests that  structural and kinematical nonhomology in elliptical galaxies can be  closely associated with the difference in  star formation history between elliptical galaxies formed by  dissipative galaxy merging. This result implies  that if a specific relation between the rapidity of star formation ($C_{\\rm SF}$) and galactic luminosity ($L$) is satisfied like as $C_{\\rm SF} \\propto L^{0.39}$, the required dependence of mass-to-light ratio on galactic mass in the FP  (the slope of the FP) can be explained  at least qualitatively in the present merger model. Further extensive observational studies  probably  assess the validity of the expected $L$ dependence of the $C_{\\rm SF}$ and thus determine whether or not the star formation history of merger progenitor galaxies can be a crucial factor for the creation of the FP. Although the present merger model is indeed rather idealized and less realistic for $real$  dissipative galaxy mergers with star formation  at higher redshift universe, it appears to have succeeded in demonstrating  that the origin of the FP can reflect a close physical relation between galactic luminosity and star formation history of merger progenitor galaxies, in particular, the rapidity of star formation of galaxies. What we should emphasize here is that the proposed idea is  only one of several promising ones in which physical meanings of the FP are extensively examined in variously different ways. Probably, both the slope and the tightness of the FP contain more profound meanings on the formation and evolution of elliptical galaxies than we considered in the present study. Thus, it is  our future study to  explore another  promising candidates related to  the  origin of the FP, such as the  age and metallicity difference of elliptical galaxies and  hierarchy of galaxy merging."
    },
    "9804/astro-ph9804296_arXiv.txt": {
        "abstract": "Gamma-ray bursts are characterized by a duration of milliseconds to several minutes in which an enormous amount of radiation is emitted.  The origin of these phenomena is still unknown because proposed models fail to explain all the observed features.  Our proposed solution to this conundrum is a new class of mass-exchanging binaries in which a neutron star transfers mass to a black hole.  According to recent studies binaries which contain a neutron star and a black hole are much more frequent than was previously believed.  Mass exchange is driven by the emission of gravitational waves but the redistribution of mass in the binary system prevents coalescence.  The phase of mass transfer is surprisingly stable and lasts for several thousands of orbital revolutions (about a minute).  With a simple analytic model we demonstrated that this new class of binaries could provide an excellent candidate for the observed phenomena known as gamma-ray bursts. ",
        "introduction": "The identification of the optical counterparts to gamma-ray bursts GRB~970228\\,(\\cite{GGv97}) and GRB~970508\\, (\\cite{He97}) and the measurement of a redshift of $z \\geq 0.853$ (\\cite{MDK97}) for the latter makes a cosmological origin for this phenomenon hard to circumvent. If indeed cosmological, they must be tremendously energetic: the maximum luminosity is many times the energy output of an entire galaxy and the total energy emitted is at least $10^{51}$ erg.  The rapid rise in luminosity and the short timescale variability suggest that the radiation is generated in an area of only a couple of hundred kilometers (a few light milliseconds).  The duration of the burst (seconds to minutes) indicates that something within this region is relatively stable. The complex temporal structure of the energy release reflects the activity of a highly variable inner engine (Sari \\& Piran 1997; Kobayashi et al.\\ 1997).  Suggestions have ranged from coalescing neutron-star binaries\\, (\\cite{BNP84}), compact objects merging with the central massive black hole of a galaxy\\, (\\cite{PFT94}) to hypernovae\\, (\\cite{P98}). However all these models have great difficulty explaining the duration (\\cite{M97}) and the intrinsic variability (Sari \\& Piran 1997) of the burst.  Detailed studies of the stability of binaries where one neutron star transfers mass to another neutron star reveal that coalescence is inevitable within a few orbital revolutions owing the Darwin-Riemann instability (\\cite{CE77}; \\cite{LRS93}).  Stable mass transfer can be achieved only if the mass of the secondary is smaller than 67\\% of the primary\\, (\\cite{JK92}).  A neutron star which transfers mass to a black hole with mass $\\apgt 6$\\,\\msun\\ cannot have a stable orbit\\, (\\cite{LS76}).  Binaries of a neutron star and a lower mass black hole have not been considered because a lack of observational evidence.  Recent understanding of accretion on to a neutron star (\\cite{C93}; \\cite{B95}; \\cite{FBH96a}) has changed this view completely, at least in the theoretical side. The birthrate (and also the merger rate) of such systems might exceed the number of binary pulsars by an order of magnitude (\\cite{PZY98}; \\cite{BB98}; Lipunov 1998). ",
        "conclusions": "The recent understanding of the super Eddington accretion process on to a neutron star in a common-envelope phase predicts the formation of a large number of close binaries in which a neutron star is accompanied by a low mass ($\\aplt 7$\\,\\msun) black hole. Model computations predict that the formation rate of such binaries exceeds the formation of classical high-mass binary pulsars by about an order of magnitude. These binaries are therefore expected to contribute appreciably to the rate of any associated observable phenomena.  There are three possible forms of mass transfer from a neutron star to its companion: unstable mass transfer to another neutron star and stable and unstable transfer to a black hole.  Also for gamma-ray bursts there is evidence for three types: short and faint bursts, long and bright and a third class of intermediate bursts\\,(\\cite{MFB98}).  In the stable binaries the duration of mass transfer lasts for about a minute after which the neutron star explodes.  An accretion disc with a mass $\\aplt 0.4$\\,\\msun\\ around the black hole develops about a second after the onset. The gamma-ray burst is powered by using the rotational energy of the black hole via the Blandford-Znajek (1977) mechanism. If the neutron star finally explodes it is accompanied by a burst of luminosity as well as anti neutrinos. The accompanying black hole is released with a low velocity.  As the binary spirals outwards instead of inwards a reversed chirp in the gravitational wave signal is expected, the signal-to-noise, however, drops considerably during this process.  In our model gamma-ray bursts are another exciting class of mass-exchanging binary stars.  We propose to call them gamma-ray binaries."
    },
    "9804/astro-ph9804282_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\noindent We set up a robust method to investigate the reliability of absorption features detections in LECS data of bright blazars. The method has been applied to 7 datasets. The analysis shows that similar features are common to all datasets (see next for 3C~273 SVP). This suggests that most of the features observed from 0.2 to 1 keV are possibly due to miscalibration of the instrument response, as of present date. The large deviations below 0.5 are probably due to the Carbon edge (at 0.29 keV), which can affect the residuals up to 1 keV. Of course a different interpretation of our results is still possible, namely that the common feature is indeed real. Should be that, interesting problems in the physical interpretation of blazar sources would be raised.  As discussed above, only 3C~273 SVP still shows a clear absorption feature after correction. The fact that the correction does not improve the $\\chi^2$ when fits are performed to the same model means that the feature detected in this dataset is \"unique\", in the sense that is not shared by any of the other datasets considered. One may argue that, if the \"common\" feature is indeed a miscalibration product, it should be apparent also in 3C~273 SVP. We note however  that the \"unique\" absorption feature is probably deep enough to overwhelm the \"common\" feature."
    },
    "9804/astro-ph9804141_arXiv.txt": {
        "abstract": "We present constraints on the evolution of large-scale structure from a catalog of 710,000 galaxies with $I_{AB} \\le 24$ derived from a KPNO 4m CCD imaging survey of a contiguous $4^{\\circ} \\times 4^{\\circ}$ region. The advantage of using large contiguous surveys for measuring clustering properties on even modest angular scales is substantial: the effects of cosmic scatter are strongly suppressed. We provide highly accurate measurements of the two-point angular correlation function, $\\omega(\\theta)$, as a function of magnitude on scales up to $1.5^{\\circ}$. The amplitude of $\\omega(\\theta)$ declines by a factor of $\\sim 10$ over the range $16 \\le I \\le 20$ but only by a factor of $2 - 3$ over the range $20 < I \\le 23$. For a redshift dependence of the spatial correlation function, $\\xi(r)$, parameterized as $\\xi(r,z) = ({r \\over r_o})^{-\\gamma} (1 + z)^{-(3+\\epsilon)}$, we find $r_o = 5.2\\pm0.4 h^{-1}$ Mpc, and $\\epsilon \\gsim 0$ for $I \\le 20$. This is in good agreement with the results from local redshift surveys. At $I > 20$, our best fit values shift towards lower $r_o$ and more negative $\\epsilon$. A strong covariance between $r_o$ and $\\epsilon$ prevent us from rejecting $\\epsilon > 0$ even at faint magnitudes but if $\\epsilon > 1$, we strongly reject $r_o \\lsim 4h^{-1}$ Mpc (co-moving). The above expression for $\\xi(r,z)$ and our data give a correlation length of $r_o(z = 0.5) \\approx 3.0\\pm0.4h^{-1}$ Mpc, about a factor of 2 larger than the correlation length at $z = 0.5$ derived from the Canada--France Redshift Survey (CFRS; \\cite{lefev96}). The small volume sampled by the CFRS and other deep redshift probes, however, make these spatial surveys strongly susceptible to cosmic scatter and will tend to bias their derived correlation lengths low. Our results are consistent with redshift distributions in which $\\sim30-50$\\% of the galaxies at $I = 23$ lie at $z > 1$. The best fit power law slope of the correlation function remains independent of $I$ magnitude for $I \\le 22$. At fainter limits, there is a suggestive trend towards flatter slopes that occurs at fluxes consistent with similar trends seen by \\cite{nwind} and \\cite{campos}. The galaxy counts span 11 magnitudes and provide an accurate calibration of the galaxy surface density. We find evidence for mild galaxy evolution -- about 1 mag of brightening or a doubling of the density by $I=23$ relative to an $\\Omega_o = 1$ no evolution model; about 0.5 mag of brightening or a factor of 1.5 increase in surface density relative to an open model. Our galaxy counts agree well with those from the HDF survey and, thus, argue against a significant inclusion of sub-galactic components in the latter census for $I < 24$. ",
        "introduction": "The evolution of large-scale structure (LSS) in the universe probes the abundance and form of dark matter, the mean baryon density, the turnover scale in the perturbation power spectrum, and the formation processes of galaxies and clusters. Observational constraints on the evolution of LSS provide strong limits on structure formation theories because the sensitivity of the mass function to $\\Omega_{\\circ}$ is high and because coherent structures on scales $>10h^{-1}$ Mpc are still in the linear regime, precisely where important distinctions can be made between competing theories. A study of this evolution is, thus, fundamental. Substantial progress has been made in studying LSS evolution from deep redshift surveys. The general consensus from the most distant probes (\\cite{cohen}, \\cite{steidel}, \\cite{mauro}) suggests that there has been significant evolution in the clustering properties of galaxies since $z \\sim 3$. Precisely just how much evolution has occurred since $z \\sim 1$ is less certain but the data are being accumulated rapidly (\\cite{CFRS}, \\cite{lefev}, \\cite{connolly}, \\cite{ajc98}).  Two-dimensional surveys offer some key advantages over narrow, deep redshift surveys like those above. Specifically, a large area angular survey subtends many galaxy correlation lengths in the transverse direction enabling many independent cells to contribute to the signal at a given depth. In addition, the power spectra derived from angular surveys are unaffected by redshift distortions, which tend to wash out small-scale power. The disadvantage of a 2D survey, of course, is that the depth sampled in a given flux range is usually broad and one must therefore deal with projection effects. Observationally, mapping large areas of sky is now quite an efficient process due to the availability of large-format CCD arrays and mosaic cameras. Ideally, one would like to conduct deep redshift surveys that also cover large, quasi-contiguous areas. As an initial step towards this goal, we have surveyed a contiguous 16 square degree area in the $I$-band using the prime focus CCD camera on the Mayall 4m at Kitt Peak National Observatory. The area is about an order of magnitude larger than any comparably deep, contiguous survey yet published. Our goal is to set accurate limits at $z = 1$ on the amplitude of structures now seen at $z = 0$.  The study of large-scale structure has advanced primarily through the study of the moments of the galaxy distribution. These fundamental functions, which must be explained by any viable structure formation model, are determined, in part, by the initial conditions at recombination and by subsequent non-linear growth. The presence of non-zero high order moments ($n \\ge 3$) provide constraints on the non-Gaussian properties of the galaxy distribution and, hence, provide very strong discrimination between competing models. The first two moments, however, are essential for assessing departures from a non-evolving, Poisson distribution. The combination of area and depth provided by this survey ($\\sim 710,000$ galaxies with $I_{AB} \\le 24$) allows us to determine the first two moments with significantly lower uncertainties than previous works. For example, the angular two-point correlation function at faint magnitudes has been measured on scales of $\\lsim 10$ arcminutes (\\eg \\cite{ef91}, \\cite{nwind}, \\cite{campos}, \\cite{lp96}, \\cite{wf97}, \\cite{bs97}) but there is substantial scatter in the results at $I \\gsim 21$, which has lead to conflicting conclusions about the evolution of structure. A significant reason for the scatter is the relatively small ($\\lsim 0.25$ deg$^2$) contiguous areas used. At these scales, one is subject to both cosmic variations in clustering and systematic corrections that are comparable in amplitude with the signal being measured. Such effects are negligible for the present survey and thus enable a major enhancement to the precision and accuracy of LSS and evolution measurements. Furthermore, the survey supports a reliable determination of at least the 3rd and 4th moments of the galaxy distribution.  In this paper we present the $I$-band galaxy counts and constraints on the evolution of the two-point correlation function. An analysis of the higher moments ($n = 3,4$) of the galaxy distribution will be presented in a separate paper (\\cite{npt98}). Section 2 contains a brief description of the observations and survey strategy. Section 3 presents results for the $I$-band galaxy counts. Our results for the two-point angular correlation function are presented in Section 4. The interpretation of our results is discussed in Section 5. We adopt $h = H_o/$(100 km sec$^{-1}$ Mpc$^{-1}$). ",
        "conclusions": "Our constraints on $\\omega(\\theta)$ reveal a dependence of its amplitude on median $I$ magnitude that appears to fit the model of the redshift evolution of $\\xi(r)$ parameterized as $\\xi(r,z) = ({r \\over r_o})^{-\\gamma} (1 + z)^{-(3+\\epsilon)}$. In fact, this model provides an acceptable representation of the data over the range $16 < I < 25$  if the results of \\cite{bs97} are also included in the fits. The decline in the amplitude of the correlation function (over a wide range of scales) with $I$ magnitude is about 3 times steeper over the range $16 \\lsim I \\lsim 19$ than over the range $20 \\lsim I \\lsim 23$. The $N(z)$ distributions that yield good fits to magnitude-$\\omega(\\theta)$ relation typically have about 40\\% of the galaxies at $z > 1$ by $I = 23$, consistent with the conclusions of \\cite{bs97}. A shallower decline in the amplitude of $\\omega(\\theta)$ with magnitude for $I \\gsim 21$ can be induced by the reduction of the proper volume element at higher $z$. \\cite{bs97} use this argument to support the assertion that a significant fraction of faint galaxies are at $z \\gsim 1$. However, a high-$z$ cutoff can also induce similar behavior since in this case the effective depth and redshift distribution are not changing dramatically as one goes to fainter limits. If such a cutoff were introduced, as noted by \\cite{lp96}, by a selection bias {\\it against} the detection of high-$z$ galaxies a shallow decline in correlation length with magnitude would result. However, a strong redshift cutoff would also effect the number counts and the corresponding turnover in the counts is not seen (see also \\cite{wf97}). Given the excellent agreement we find with the HDF number counts, we conclude that a selection bias of this nature is not a significant problem for our sample.  Taken at face value, our best fits to the above model suggest the evolution of the spatial correlation function is relatively mild over the depth of the survey. When we allow $r_o$ to be a free parameter in the fit, we obtain exponents in the range $-1.7 \\lsim \\epsilon \\lsim -0.2$ with $3.8 \\lsim r_o \\lsim 5.6h^{-1}$ Mpc depending on the assumed redshift distribution. Growth in the clustering of galaxies (in proper coordinates) requires $\\epsilon > 0$. The 95\\% confidence limits are substantial, nevertheless, given the strong covariance between $r_o$ and $\\epsilon$. As Figure~\\ref{chisqr} demonstrates, our data are also consistent with $\\epsilon > 0$ providing $4.5 \\lsim r_o \\lsim 7h^{-1}$ Mpc. If $\\epsilon \\gsim 1$ then values of $r_o \\lsim 4h^{-1}$ Mpc are strongly rejected. As already noted, redshift surveys of local galaxies find $r_o \\sim 5.5\\pm1.5h^{-1}$ Mpc.  The interpretation of the goodness of the fit to the above model, however, must take into account a number of additional selection effects that may mask more substantial evolution of the correlation function. It is known that low-$z$ galaxy clustering depends on both the morphology and luminosity of the objects being studied ({\\it e.g.}, \\cite{dg76}, \\cite{moore94}, \\cite{love95}, \\cite{guzzo97}, \\cite{vl97}). The sense of these trends is that elliptical galaxies and more luminous galaxies tend to have larger correlation lengths than spiral galaxies and less luminous galaxies, respectively. We also know that the mean absolute luminosity of the objects in a flux-limited sample will increase with redshift. The relative insensitivity of our derived $r_o$ values to $I$ magnitude for galaxies with $I \\ge 17$ could therefore be, in part, the result of two competing effects. As one goes fainter the mean redshift of the galaxies in the survey tends to increase. If galaxies at $0.5 \\lsim z \\lsim 1$ exhibit a similar luminosity dependent clustering as local galaxies and/or $I-$band selection enhances the elliptical galaxy fraction at these redshifts then it is possible that a decrease in the spatial correlation length with magnitude could be partially offset by the tendency for more luminous and/or early type galaxies to be more strongly clustered. This could also explain the observed preference for negative $\\epsilon$ values (which suggests only modest clustering evolution). The precise evolution of $\\epsilon$ and $r_o$ with redshift is intimately connected to the density pertubation power spectrum and the galaxy merger rate (\\cite{mosc98}). The complexity of the dependence makes a unique theoretical interpretation the observational constraints problematic. None the less, the constraints are now solidly determined on degree-scales and less down to $I = 23$ from the present survey.  More stringent constraints on clustering evolution will require the addition of multiple passbands or redshift data. \\cite{ef91} report a marginal increase in the clustering amplitude as one selects redder passbands and posit that normal galaxies dominate the composition of fainter $I-$band selected galaxy samples as opposed to the weakly clustered faint galaxies, that appear to dominate $U$ and $B$ selected surveys. \\cite{lp96} also find a marginally significant color dependence, showing that red ($V-I > 1.5$) galaxies have a larger clustering amplitude than blue galaxies. \\cite{lefev96} find no color dependence to the clustering properties of galaxies at $z > 0.5$ from an analysis of the CFRS data. They have also measured the spatial correlation function directly from their redshift data. They find a correlation length of $r_o = 1.57\\pm0.09h^{-1}$ Mpc  at $z \\sim 0.5$ which, if they are sampling a mix of galaxies similar to those in local redshift surveys, implies $0 < \\epsilon < 2$. Our results are consistent with this $\\epsilon$ range. From equation (2), we know $r_o(z) = r_o(0) \\left( 1 + z \\right)^{(-(3+\\epsilon)/\\gamma)}$. Using our best-fit results for the CFRS redshift distribution we, thus, find that $r(z=0.5) \\approx 3.0h^{-1}$ Mpc. Results for the EvLF redshift distribution models yield $r(z=0.5) = 2.7 \\pm 0.4h^{-1}$ Mpc. The CFRS galaxies were also $I-$band selected so any difference in results between the CFRS and our angular survey are due either to projection effects (which tend to wash out clustering) or the volume sampled as opposed to intrinsic differences between the galaxy population.  \\cite{ajc98} determined the redshift dependence of the amplitude of $\\omega(\\theta)$ using photometric redshifts derived from the HDF survey. For $z > 0.4$, they find $r_o = 2.37h^{-1}$ Mpc and $\\epsilon = -0.4^{+0.37}_{-0.65}$ However, in order to be consistent with low-$z$ surveys, they would require $\\epsilon \\sim 2$. They suggest, therefore, that the expression for the evolution of the spatial correlation function, $\\xi(r,z) = ({r \\over r_o})^{-\\gamma} (1 + z)^{-(3+\\epsilon)}$, is not particularly good. Others have made similar claims based on poor fits to the $\\omega(\\theta)$ vs. magnitude relation. This is in contrast with the excellent fit we obtain extending over $7$ magnitudes (Figure~\\ref{wonefits}). We speculate that the discrepancy is caused by cosmic variance due to the small area used in previous measurements of the correlation function. \\cite{ajc98} note that the HDF survey subtends only $800h^{-1}$ kpc at $z = 1$ (q$_o = 0.1$), a fraction of the galaxy correlation length. The CFRS is larger (five $10' \\times 10'$ fields) but still covers a relatively small volume ($10'$ spans $3.1h^{-1}$ Mpc at $z = 1$).  Hence, these surveys are also subject to the same sorts of problems associated with cosmic scatter and undersampling of LSS as small area angular surveys. Indeed, \\cite{delap} make a clear demonstration of this point at low-$z$: even in their redshift survey of $\\sim 1810$ galaxies ($z \\le 0.05$) filling a volume of $5 \\times 10^5h^{-3}$ Mpc$^3$, the density fluctuations caused by LSS prevent the determination of $r_o$ by better than a factor of 2. Given that the deep, redshift probes sample roughly comparable volumes (but with the added complication that they stretch over a significant fraction of cosmic time), it seems optimistic to expect reliable constraints to be determined from the existing surveys. In fact, it is quite remarkable that the values obtained from them are in agreement with the local correlation length to within a factor of two. Not until much larger volumes are surveyed will the constraints on $r_o$ at high redshift be robustly measured. If the results from direct measurements of the spatial correlation function at higher redshifts are ultimately accurate, however, and assuming their sample composition is not significantly different from the low redshift surveys, then the basic model may be in need of refinement.  One possible refinement, discussed in some detail already by \\cite{nwind}, is to include a redshift-dependent spatial correlation function slope. This would be a natural consequence of models in which there is a scale dependence to the growth of structure. For example, if galactic scale structures grow significantly faster than structures on scales $\\gsim 20h^{-1}$ Mpc the correlation function slope would steepen with time. This would also cause the slope of $\\omega(\\theta)$ to flatten at fainter magnitudes on small scales and, depending on the details of the density perturbation power spectrum and the merger rate, could also result in an increase in $\\epsilon$ with redshift on small scales. Adopting the parameterization $\\gamma(z)=1.8(1+z_{med})^{-C}$ proposed by \\cite{nwind}, implies $C = 0.35\\pm0.10$ which is consistent with their limit. As \\cite{nwind} demonstrate, values in the range $C \\lsim 0.4$, are consistent with structure formation models which have significant power on co-moving scales of $\\lsim 10h^{-1}$ Mpc but disfavor models with little small scale power such as HDM. However, since the flattening is only seen at $I \\gsim 22$, the above parameterization of the slope redshift dependence is not particularly good -- there appears to be no real redshift dependence to the slope until $z_{med} \\gsim 0.6$."
    },
    "9804/astro-ph9804188_arXiv.txt": {
        "abstract": "We present a new determination of the mass function of galaxy clusters, based on optical virial mass estimates for a large sample of 152 nearby ($z\\le 0.15$) Abell--ACO clusters, as provided by Girardi et al. (1998). This sample includes both data from the literature and the new ENACS data. The resulting mass function is reliably estimated for masses larger than $M_{lim}\\simeq 4\\times 10^{14}$\\msun, while it is affected by sample incompleteness at smaller masses.  We find $N(>M_{lim})=(6.3\\pm 1.2)\\,10^{-6}$ $(h^{-1} Mpc)^{-3}$ for cluster masses estimated within a 1.5 \\h~radius.  Our mass function is intermediate between the two previous estimates by Bahcall \\& Cen (1993) and by Biviano et al. (1993).  Based on the Press--Schechter approach, we use this mass function to constrain the amplitude of the fluctuation power spectrum at the cluster scale.  After suitably convolving the PS predictions with observational errors on cluster masses and {\\sl COBE}--normalizing the fluctuation power spectrum, we find $\\sigma_8=(0.60\\pm 0.04) \\Omega_0^{-0.46+0.09\\Omega_0}$ for flat low--density models and $\\sigma_8=(0.60\\pm 0.04) \\Omega_0^{-0.48+0.17\\Omega_0}$ for open models (at the 90\\% c.l.).    \\vspace*{6pt} \\noindent {\\em Subject headings: } galaxies: clusters: general - cosmology: observations - cosmology: theory - large scale structure of universe. ",
        "introduction": "According to the scenario of hierarchical formation of cosmic structures, clusters of galaxies arise from rare high peaks of the initial density fluctuation field and, at present time, represent the largest virialized cosmic structures. As a consequence, standard analytical arguments based on the Press \\& Schechter (1974; hereafter PS) approach show that their abundance is highly sensitive to the amplitude of the density fluctuations on the cluster mass scale (e.g., White, Efstathiou \\& Frenk 1993). In particular, if $M$ is the typical cluster mass and $\\Omega_0$ is the cosmic density parameter, then the typical size of the density fluctuation collapsing into a cluster is $R\\propto (\\Omega_0 M)^{1/3}$. Typical cluster masses, $M\\sim 5\\times 10^{14}$\\msun \\footnote{Here $h$ is the Hubble constant in units of 100$\\vel$Mpc$^{-1}$}, are rather close to the average mass contained within a sphere of 8\\h~ radius. Therefore, it is common to express constraints from the observed local cluster mass function (i.e., the number density of clusters of a given mass), $n(M)$, in terms of $\\Omega_0$ and $\\sigma_8$, the r.m.s. density fluctuation within a sphere of 8\\h~ radius (e.g., Eke, Cole, \\& Frenk 1996; Viana \\& Liddle 1996).  From the theoretical side, the mass function of clusters is easily obtainable either from N--body simulations or, to a comparable degree of reliability (e.g., Lacey \\& Cole 1993; Eke et al. 1996; Borgani et al. 1998), by applying the analytical PS recipe. On the contrary, predictions of cosmological models about quantities which are more directly connected to observations, like the internal velocity dispersion of member galaxies, $\\sigma_v$, and the $X$--ray temperature of the intracluster gas, $T_X$, are much less straightforward and rely on assumptions which are quite debated. As for $\\sigma_v$, N--body simulations allow reliable predictions of the internal velocity dispersion of clusters, but only for the DM particles (e.g., Frenk et al. 1990; Borgani et al. 1997; van Haarlem, Frenk, \\& White 1997), while the connection to the observed $\\sigma_v$ requires a plausible model of galaxy formation. Numerical simulations by Frenk et al. (1996) have shown that present uncertainties about the physics of galaxy formation in clusters can lead to an incorrect determination of the virial mass by up to a factor five.  As for $T_X$, it is generally believed to be a rather accurate indicator of the cluster mass (e.g., Evrard, Metzler, \\& Navarro 1996). On the other hand, both analytical predictions and hydrodynamical cluster simulations, based only on the adiabatic physics of the intracluster gas, are able to reproduce neither the observed relation between bolometric $X$--ray luminosity and temperature (e.g., Eke, Navarro \\& Frenk 1997; Bryan \\& Norman 1998, and references therein), nor the observed size of core radii in the gas distribution (e.g., Makino, Sasaki, \\& Suto 1998, and references therein).  Such a discrepancy creates a need to introduce additional physics to explain the thermodynamics of the the intra--cluster gas (e.g., Bower 1997; Cavaliere, Menci, \\& Tozzi 1997, and references therein).  From the observational side, cluster masses are inferred from either $X$--ray or optical data, under the general hypothesis of dynamical equilibrium.  Estimates based on gravitational lensing do not require assumptions about the dynamical status of the cluster, but a good knowledge of the geometry of the potential well is necessary (e.g., Narayan \\& Bartelmann 1997).  Claims for a discrepancy (by a factor 2-3) between cluster masses obtained with different methods casted doubts about the general reliability of mass estimates (e.g., Wu \\& Fang 1997).  However, recent analyses have shown that such discrepancies can be explained by the different way in which strong cluster substructures bias mass estimates based on different methods (Allen 1997; Girardi et al.  1997). In particular, the analysis by Girardi et al. (1998, hereafter G98) demonstrated that, for nearby clusters without strong substructures, a good overall agreement exists between $X$--ray and optical virial mass estimates.  A reliable determination of cluster masses is a necessary but not sufficient condition to guarantee a reliable estimate of the mass function. Indeed, the latter also requires a cluster sample which is large enough to be a fair representation of the cluster population in the local Universe, at least above a certain mass limit. In past years, such requirements limited the possibility of obtaining a stable result about the cluster mass function and, in fact, the only two estimates so far presented (Bahcall \\& Cen 1993, hereafter BC93; Biviano et al. 1993, hereafter B93) turn out to be quite discrepant. BC93 derived cluster masses by using global scaling relations to connect mass to cluster richness and $X$--ray temperature. B93 estimated cluster masses from a homogeneous virial analysis of about 70 clusters and found a larger abundance, by about one order of magnitude, at the mass scale $M\\simeq 5\\times 10^{14}$\\msun (cf. also Figure~3 below).  A new and more robust determination of the observational mass function based on optical virial analysis is now possible thanks to the recent availability of a large amount of redshift data for rich clusters, as provided by the ESO Nearby Abell Cluster Survey (ENACS; Katgert et al. 1996; 1998).  Recently, G98 computed virial masses in a homogeneous way for a large sample of $\\sim 170$ clusters, obtained by combining the ENACS sample with data from other authors (see also Fadda et al. 1996; hereafter F96). The extension of this sample makes it the largest available data set over which to reliably estimate the cluster mass function.  In the following analysis we will use the G98 results with the twofold aim of {\\em (a)} estimating the optical virial mass function of nearby galaxy clusters, and {\\em (b)} investigating the subsequent constraints on the power spectrum of cosmic density fluctuations.  This paper is organized as follows. In \\S~2 we briefly describe the statistical sample of clusters masses as obtained by G98 and then present the determination of the cluster mass function. Afterwards, we compare it with previous determinations based on optical data by BC93 and B93. In \\S~3 we apply the PS approach to place constraints on the fluctuation power spectrum for different cosmological models. We give a brief summary of our main results and draw our conclusions in \\S~4. ",
        "conclusions": ""
    },
    "9804/astro-ph9804231_arXiv.txt": {
        "abstract": "We have made the first detection of circumstellar absorption lines of the $^{12}$C$^{13}$C \\phsys~ (Phillips) system  1-0 band and     the $^{12}$C$^{16}$O \\cosys~ first-overtone     2-0 band in the spectrum of the post-AGB star HD\\,56126 (IRAS\\,07134+1005). All current detections of circumstellar molecular absorption lines towards HD\\,56126 ($^{12}$C$_{2}$, $^{12}$C$^{13}$C, $^{12}$C$^{14}$N, $^{13}$C$^{14}$N, and $^{12}$C$^{16}$O) yield the same heliocentric velocity of $v_{\\rm CSE}=77.6\\pm0.4$~km~s$^{-1}$. The $^{12}$C$_{2}$, $^{12}$C$^{13}$C, and $^{12}$C$^{16}$O lines give rotational temperatures and integrated column densities of $T_{\\rm rot} = 328\\pm37$~K, $\\log N_{\\rm int}=15.34\\pm0.10$~cm$^{-2}$, $T_{\\rm rot} = 256\\pm30$~K, $\\log N_{\\rm int}=13.79\\pm0.12$~cm$^{-2}$, and $T_{\\rm rot} =  51\\pm37$~K, $\\log N_{\\rm int}=18.12\\pm0.13$~cm$^{-2}$ respectively. The rotational temperatures are lower for molecules with a higher permanent dipole moment. Derived relative column densities ratios are $^{12}$C$_{2}$/$^{12}$C$^{13}$C$=36\\pm13$, and $^{12}$C$^{16}$O/($^{12}$C$_{2}$+$^{12}$C$^{13}$C)$=606\\pm230$. Combined with data from Paper~III we find relative column densities of $^{12}$C$^{16}$O/($^{12}$C$^{14}$N+$^{13}$C$^{14}$N)$=475\\pm175$ and $^{12}$C$^{14}$N/$^{13}$C$^{14}$N$=38\\pm2$.  Under chemical equilibrium conditions, $^{12}$C$^{13}$C is formed twice as easily as $^{12}$C$_{2}$. The isotopic exchange reaction for $^{12}$C$_{2}$ is too slow to significantly alter the $^{12}$C$_{2}$/$^{12}$C$^{13}$C ratio and the $^{12}$C$_{2}$ to $^{12}$C$^{13}$C ratio a good measure of half the carbon isotope ratio: $^{12}$C/$^{13}$C=$2\\times$$^{12}$C$_{2}$/$^{12}$C$^{13}$C=$72\\pm26$. This is in agreeement with our prediction that the isotopic exchange reaction for $^{12}$C$^{14}$N  is efficient and our observation in Paper~III of $^{12}$C$^{14}$N/$^{13}$C$^{14}$N$=38\\pm2$.  A fit of the C$_{2}$ excitation model of van Dishoeck \\& Black (\\cite{vandishoeckblackmodel}) to the relative population distribution of C$_{2}$ yields $n_{\\rm c} \\sigma / I=3.3 \\pm 1.0 \\times 10^{-14}$. At $r\\simeq 10^{16}$~cm this translates in $n_{\\rm c}=1.7 \\times 10^{7}$~cm$^{-3}$ and $\\dot M \\simeq 2.5 \\times 10^{-4}$~M$_{\\odot}$~yr$^{-1}$.   \\noindent ",
        "introduction": "HD\\,56126  (IRAS\\,07134+1005) is in the post-AGB stage (also referred to as the Pre-Planetary Nebulae stage) of stellar evolution. This is the relative short transition stage from the Asymptotic Giant Branch (AGB) to the White Dwarf (WD) phase (Iben~\\cite{iben}). Low-mass stars (0.8 M$_{\\odot} \\leq$ M$_{\\ast} \\leq$ 8.0 M$_{\\odot}$) may experience up to three dredge-up episodes. The first occurs on the Red Giant Branch (RGB). The second on the Early-AGB (E-AGB) is limited to stars with initial masses in the range of 4 to 8~M$_{\\odot}$. The third occurs on the Thermal Pulsating-AGB (TP-AGB). The last dredge-up changes the surface abundances of the star most drastically: most noticeable the carbon and s-process elements abundances enhanced (Forestini \\& Charbonnel~\\cite{forestinicharbonnel}). Several studies of the photospheric abundance of post-AGB stars have indeed shown that some of these stars are carbon rich and have enhanced abundances of s-process elements (Luck \\& Bond~\\cite{luckbond};  % Klochkova~\\cite{klochkova};    % van Winckel~\\cite{vanwinckel}; % Reddy et al.~\\cite{reddyetal},\\cite{reddybakkerhrivnak}; % Decin et al.~\\cite{denin}). % Since AGB stars lose mass through a dense stellar wind, post-AGB stars are surrounded by a shell of material which expands at  typically 3 to 30~km~s$^{-1}$. This material constitutes the circumstellar environment (CSE) of the star. The CSE material mixes with interstellar clouds and forms the parent material for the next generation of stars. In this manner, nucleosynthesis products are used to build  new  stars and the Galaxy is enriched with heavy elements (elements more massive than helium). The study of the chemical structure of the CSE of evolved stars is therefore of importance to understand the chemical enrichment of the Universe.  Stellar evolution theory needs constraints from observations. The  $^{12}$C/$^{13}$C suits this goal very well since during the AGB evolution $^{12}$C is formed in the helium burning layer, while $^{13}$C is formed as a by-product in the hydrogen burning layer. The observed ratio of these two isotopes depends on the efficiency and the relative strength of these two layers, and on the efficiency of the dredge-up processes.  Theoretical models of the chemical structure of the CSE of highly evolved stars have mainly concentrated on understanding the well-studied, and well-observed carbon star IRC\\,+10216 (for references see Olofsson \\cite{olofsson}). When a star reaches the tip of the AGB, the mass-loss rate has its maximum strength and drops several orders of magnitude after the star leaves the AGB. Material ejected during the AGB phase slowly expands from the star and leaves a cavity in the CSE which will grow with time. Eventually leading to a cavity the size of a planetary nebulae. The elemental composition of the CSE is fixed after the material leaves the surface of the star. In contrast, the molecular composition of the CSE is constantly changing under the influence of the stellar and interstellar radiation fields: molecules are destroyed and formed continuously. Most studies of the CSE make use of molecular line emission in the sub-millimeter and radio (e.g. CO and HCN) while recently the ISO satellite has allowed a study of molecular lines in the infrared. Emission lines of spatially unresolved sources provide no positional information, and only partial velocity information (the radial component). This is obviously not the case for absorption lines since they are formed in a pencil-beam towards the continuum source and probe therefore a well defined region of the CSE. Secondly, a molecular absorption line spectrum has many lines allowing a study of the population of the energy levels of the molecules and  processes like optical pumping. Absorption lines therefore yield information additional to that already available from emission line studies. As part of an ongoing effort to study the characteristics of post-AGB stars, we have been studying circumstellar $^{12}$C$^{14}$N and $^{12}$C$_{2}$ lines in the optical spectra of post-AGB stars (referred to as C2CN stars) which show the unidentified 21~$\\mu$m feature (Kwok et al. \\cite{kwok89},\\cite{kwok95}; Hrivnak \\& Kwok \\cite{hrivnakkwok}; Justtanont et al. \\cite{justtanont}). These electronic bands allow an accurate determination of the expansion velocity of the CSE, rotational temperature, column density, and molecular column density ratios (Bakker et al. \\cite{paperI} (Paper~I), \\cite{paperII} (Paper~II), Bakker \\& Lambert \\cite{paperIII} (Paper~III)).  In Paper~III we presented measurements of the ratio CN/$^{13}$CN (from now on C means $^{12}$C, N means $^{14}$N, and O means $^{16}$O). We found a ratio of $38\\pm2$ but argued that the true C/$^{13}$C ratio is likely closer to $67$ owing to the isotopic exchange reaction ($ {\\rm ^{13}C^+} + {\\rm      CN}   \\rightleftharpoons {\\rm      C^+} + {\\rm ^{13}CN} + \\Delta E $). In this paper (IV) we present measurements of the C$_{2}$/C$^{13}$C ratio. Combining these two independent isotope ratios allows us to constrain the isotope exchange reaction and give a more accurate determination of the true C/$^{13}$C ratio in the CSE surrounding HD\\,56126.   In Sec.~2 we discuss the optical and infrared observations and the source of the data on the equivalent widths used in our analysis. Sec.~3 describes the  sources from which we obtained the molecular data (C$_{2}$, C$^{13}$C, and CO). Sec.~4 is the analysis and Sec.~5 the discussion. ",
        "conclusions": "The rotational temperatures of C$^{13}$C and CO are well constrained by a linear fit to their rotational curves (Fig.~\\ref{fig_rot}). Within the error C$_{2}$ can be fitted to the same rotational temperature as C$^{13}$C, but it is clear that the C$_{2}$ rotational curve  has a depression from a linear curve near $E(J^{\\prime\\prime})/k=300$~K. For lower energies (lower J$^{\\prime\\prime}$), the curve is steeper which translates in a lower rotational temperature, while for higher energies (higher J$^{\\prime\\prime}$), the curve is flatter which translates in a higher rotational temperature. This effect has been noted before (Paper~I, II) and is characteristic for a molecule which is not in LTE and has a rotational temperature in excess of the kinetic temperature of the gas. C$_{2}$ is a homo-nuclear molecule and has therefore no permanent dipole moment. Selection rules do not allow rovibrational (infrared) and rotational (sub-millimeter) transitions. The molecule absorbs energy from the infrared to the ultraviolet (the observed molecular bands), but has no strong cooling mechanism. The result is that the molecule is rotationally excited above the kinetic temperature. An overall fit to the curve yields $T_{\\rm rot}=328\\pm37$~K. C$^{13}$C has a very small electric dipole moment and very weak allowed rovibrational and rotational transitions allow the molecule to cool. The possibly lower rotational temperature for C$^{13}$C of $T_{\\rm rot}=256\\pm30$~K is consistent with this model. CO has a  permanent dipole moment ($0.1222$~Debye, for references see Kirby-Docken \\& Liu \\cite{kirbydockenliu77}). It can effectively cool by emitting a photon in the infrared or sub-millimeter. It is thus not unexpected that we find a rotational temperature lower than that for C$_{2}$ and C$^{13}$C. CN, with a permanent dipole moment of $1.45\\pm0.08$~Debye (Thomson \\& Dalby \\cite{thomsondalby}), cools more efficiently and has a lower rotational temperature of $T_{\\rm rot}=11.5\\pm0.6$~K. Besides the difference in dipole moment, there are second order factors which should be considered in order to understand the excitation of the molecule. Since each molecules has its own characteristic absorption spectrum, different molecules absorb different parts of the stellar spectrum. The energy distribution of the central stars, the star's photospheric spectrum, and the reddening between the star and the line forming region should therefore be considered. Clearly, it would be of great interest to model the excitation of the five molecules detected and extract information about the conditions within the CSE from their excitation.  Because of the difference in rotational temperature between C$_{2}$ and C$^{13}$C, and the fact that C$^{13}$C has twice as many $J^{\\prime\\prime}$ levels available, the C$_{2}$/C$^{13}$C can not be determined for each  $J^{\\prime\\prime}$ level independently. Instead the only acceptable way is by integrating the population over all available levels. The unobserved levels for C$^{13}$C contain a significant fraction of the total column density and the total C$^{13}$C column density can therefore only be obtained by computing a rotational temperature and extrapolating to the unobserved energy levels (labeled as $\\log N_{\\rm int}$ in Table~\\ref{tab_logn}). Using the integrated column densities for C$_{2}$ and C$^{13}$C we find C$_{2}$/C$^{13}$C$=36\\pm13$.  The conversion from C$_{2}$/C$^{13}$C to C/$^{13}$C is rather straightforward. Under chemical equilibrium conditions, C$^{13}$C is formed twice as easily as C$_{2}$ (Tatum \\cite{tatum}). This can be best demonstrated with Tatum's example. In a draw there are equal number of white and black socks ($^{12}$C and $^{13}$C atoms to make C$_{2}$H$_{2}$ $\\rightarrow$ C$_{2}$H $\\rightarrow$ C$_{2}$). By taking at random two socks out of the draw, the possible combinations are (white,white), (white,black), (black,white), and (black,black). Since (white,black) is the same as (black,white) the probability is 1:2:1 for a white, mixed, and black pair of socks. Hence the equilibrium abundance of the hetero-nuclear molecule must be twice that of the homo-nuclear molecule. Excluding any other reaction this leads to C/$^{13}$C$=2 \\times$C$_{2}$/C$^{13}$C=$72\\pm26$.  CN is a polar molecule with a permanent dipole moment. The rate coefficient for the isotopic exchange reaction of CN with C$^{+}$ is as high as $10^{-7}$~cm$^{-3}$~s$^{-1}$ (Adams et al. \\cite{adamsetal}). Since $^{13}$CN has a 31~K lower zero-point energy than CN, the reaction making $^{13}$CN from CN is exothermic. In paper~III we constructed a simple chemical model of the CSE of HD\\,56126 and found that for C$_{2}$/$^{13}$CN$=38\\pm2$ we get C/$^{13}$C$\\simeq 67$. C$_{2}$ is a symmetric molecule possessing no permanent dipole moment, theory predicts a rate coefficient which is typically $10^{-9}$~cm$^{-3}$~s$^{-1}$. Since C$^{13}$C has a 25~K lower zero-point energy than C$_{2}$ the reaction making C$^{13}$C from C$_{2}$ is exothermic. Since the zero-point energy difference for C$^{13}$C and $^{13}$CN to C$_{2}$ and CN respectively are about the same, and C$_{2}$ and CN have about the same abundance with distance from the star, we can use the chemical model of Paper~III (their Fig.~6.d.) to assess the importance of the isotopic exchange reaction for C$_{2}$. We find that our model predicts C$_{2}$/C$^{13}$C=60 for C/$^{13}$C=67. Taking into account the many uncertainties of this model (C$^{+}$ abundance, kinetic temperature etc.) we argue that within the errors of our computations, the isotopic exchange reaction does not significantly alter the C$_{2}$/C$^{13}$C ratio in the circumstellar shell surrounding HD\\,56126.  Given these argument for C$_{2}$ and CN, we argue that the intrinsic carbon isotope ratio is well constrained by the C$_{2}$/C$^{13}$C ratio as C/$^{13}$C=$72\\pm16$, and that a rate coefficient for the CN isotopic exchange reaction of $10^{-7}$~cm$^{-3}$~s$^{-1}$ is consistent with the observed CN/$^{13}$CN ratio.  The prototype of the AGB stars, IRC~+10216/CW Leo, is a massive, highly-evolved carbon star with $3 M_{\\odot} \\leq M_{\\rm ZAMS} \\leq 5 M_{\\odot}$ (Gu\\'{e}lin et al. \\cite{guelinetal}). For IRC~+10216 the isotope ratios are well determined and give an estimate of the ratios which one might expect to detect for a carbon-rich post-AGB star like HD~56126: C/$^{13}$C=44$^{+3}_{-3}$. (see Forestini \\& Charbonnel  \\cite{forestinicharbonnel} for  an overview). Our estimate of the C/$^{13}$C ratio for HD~56126's shell is consistent with that of IRC~+10216. Results for circumstellar shells around four other carbon stars were provided by Kahane et al. (\\cite{kahane}): C/$^{13}C \\leq 60$. These too suggest that HD~56126 is not exceptional.   Combining our results from Paper~III and this paper yields a column density ratio  of CO/(C$_{2}$+C$^{13}$C)$=606\\pm230$ and CO/(CN+$^{13}$CN)$=475\\pm175$. These column density ratios do not necessarily reflect abundance density ratios. CO is abundant throughout the whole CSE, while C$_{2}$ and CN are only present in a shell within the CSE where interstellar photons photo-dissociate HCN to CN and  C$_{2}$H$_{2}$ via C$_{2}$H to C$_{2}$ (Cherchneff et al. \\cite{cherchneff}). Column density ratios of CO to C$_{2}$ and CN (excluding their isotopes) have been measured in several ways. Olofsson (\\cite{olofsson}) lists abundances for a large range of molecules detected in the shell surrounding AGB stars (although most molecules are only detected in IRC\\,+10216). From his work we find CO/C$_{2}=1 \\times 10^{3}$ and CO/CN$=1 \\times 10^{3}$. Bachiller et al. (\\cite{bachilleretal}) present abundance ratios for, among others,  CRL~2688 (Egg Nebulae) obtained  from sub-millimeter line emission observations CO/CN=909. Within a factor of two these numbers are consistent with what we find for HD\\,56126. Bachiller et al. suggest that the CO/$^{13}$CO$\\simeq 20$ likely reflects the C/$^{13}$C for the planetary nebulae. If so, their C/$^{13}$C is a factor three lower than ours. This could suggest that the objects observed by Bachiller et al. are not the precursors of HD\\,56126 but possible PN's with higher mass progenitor ($M_{\\ast} \\geq 2$~M$_{\\odot}$) while the progenitor of HD\\,56126  has ($M_{\\ast} \\leq 2$~M$_{\\odot}$).  A model for the C$_{2}$ molecule has been constructed by van Dishoeck \\& Black (\\cite{vandishoeckblackmodel}) which takes into account collisional and radiative excitation and de-excitation. For $T_{\\rm kin}=25$~K we find $S=n_{\\rm c} \\sigma / I=3.3 \\pm 1.0 \\times 10^{-14}$~cm$^{-1}$ (see also Bakker et al. \\cite{bakkeredinburg}). It is curious that these estimates of $T_{\\rm kin}$ and $S$ are quite similar to those derived from C$_{2}$ molecules in diffuse interstellar clouds. This is surprising as in Paper~II we estimated that the circumstellar C$_{2}$ molecules are pumped largely by the stellar radiation field. From a fit to the rotational population, $\\sigma=7.8 \\times 10^{-16}$~cm$^{2}$ as the collisional cross-section for C$_{2}$-H$_{2}$ (Phillips \\cite{phillips}), and $I(r=10^{16}~{\\rm cm}) \\simeq 4.1 \\times 10^{5}$ the ratio of the stellar  radiation field  relative to the standard interstellar radiation field, we find $n_{\\rm c}=n ({\\rm H}) + n({\\rm H_{2}})=1.7 \\times 10^{7}$ as the number of collisional partners for C$_{2}$. At $r=10^{16}$~cm this translates in $\\dot M \\simeq 2.5 \\times 10^{-4}$~M$_{\\odot}$~yr$^{-1}$. This mass-loss rate as derived from modeling the C$_{2}$ excitation gives a result comparable to the mass-loss rate derived using other techniques.  Based on this work we can make some suggestions to continue this study. A critical detection would be that of the quadrupole transitions of H$_{2}$ at 2.2$\\mu$m and $^{13}$CO rovibrational bands at 4.6~$\\mu$m. This would allow to determine the column density ratio of the detected molecules relative to the most abundant specie (H$_{2}$) and the CO/$^{13}$CO ratio. Secondly, more C2CN stars should be studied to determine the C$_{2}$/C$^{13}$C and CN/C$^{13}$N ratio to get good statistics. The most ambitious project would be to model the excitation of all molecules simultaneously in order to obtain the physical conditions of the CSE (densities, radiation field, extinction etc.)."
    },
    "9804/astro-ph9804007_arXiv.txt": {
        "abstract": "We report the analysis of HST {\\tt WFPC2} F606W images of 40 spiral galaxies belonging to the sample introduced in Carollo et al. (paper I), where 35 other targets were discussed. We describe the optical morphological properties of the new 40 galaxies, derive the surface brightness profiles for 25 of them, and present the results of photometric decompositions of these profiles into a ``bulge'' ($R^{1/4}$ or exponential) and a disk component.  The analysis of the enlarged sample of 75 galaxies puts on a statistically more solid ground the main results presented in paper I, namely: {\\it I.}  In $\\approx$30\\% of the galaxies, the inner, morphologically-distinct, structures have an irregular appearance. Some of these ``irregular bulges'' are likely to be currently star forming. {\\it II.}  Resolved, central compact sources are detected in about 50\\% of the galaxies. {\\it III.}  The central compact sources in galaxies with nuclear star formation are brighter, for similar sizes, than those in non star forming galaxies. {\\it IV.}  The luminosity of the compact sources correlates with the total galactic luminosity.  \\\\ Furthermore, the analysis of the enlarged sample of 75 objects shows that: {\\it A.} Several of the non-classical inner structures are well fitted by an exponential profile. These ``exponential bulges'' are fainter than $R^{1/4}$ bulges, for given total galaxy luminosity and (catalog) Hubble type later than Sab.  {\\it B.}  Irregular/exponential bulges typically host central compact sources.  {\\it C.} The central sources are present in all types of disk galaxies, starting with systems as early as S0a. About 60\\% of Sb to Sc galaxies host a central compact source. Many of the galaxies which host compact sources contain a barred structure.  {\\it D.}  Galaxies with apparent nuclear star formation, which also host the brightest compact sources, are preferentially the early- and intermediate-type (S0a-Sb) systems. {\\it E.} None of the features depends on environment: isolated and not isolated galaxies show indistinguishable properties. \\\\ Independently from the physical nature of the non-classical inner structures, our main conclusion is that a significant fraction of galaxies classified from the ground as relatively early-type spirals shows a rich variety of central properties, and little or no morphological/photometric evidence for a smooth, $R^{1/4}$-law bulge. ",
        "introduction": "This is the second of a series of papers dedicated to the investigation of the (nuclear) properties of spiral galaxies. To this aim, we have started a HST {\\tt WFPC2} snapshot survey in the F606W filter of a sample of 92 spiral galaxies (complemented by 15 additional early-type galaxies, for a comparison with the literature), 75 of which have already been observed. In the first paper of the series (Carollo et al.\\ 1997a, paper I), we have discussed the optical nuclear morphology (and surface brightness profiles) for 35 of our target galaxies, i.e. the ones observed until the HST servicing mission of February 11, 1997. In this paper, we do the same for the additional 40 targets which have been observed since then, the last of the observations being dated April 15, 1997. Analogously to what done in paper I, we use our measurements to investigate the occurrence of nuclear spiral structure, the properties of the bulge-like components in the inner regions of disk galaxies, the frequency of occurrence of nuclear compact sources, and the relation of the above properties with global galactic properties. Here, the enlarged sample of 75 galaxies allows us to perform a statistical analysis of the frequency of occurrence of the various features as a function of luminosity, Hubble type and environment.  In contrast to the wealth of information on the nuclei of early-type galaxies provided by HST (Crane et al.\\ 1993; Jaffe et al.\\ 1994; Lauer et al.\\ 1995; Carollo et al.\\ 1997b, 1997c; Faber et al.\\ 1997), at the start of our investigation little was known either about the nuclear properties of early-type disk galaxies, or about the relation between the nuclear properties and those of the surrounding galactic structure, e.g., the bulge and/or the inner disk (see the recent review by Wyse, Gilmore \\& Franx 1997). Phillips et al.\\ (1996) presented WF/PC images for a sample of 20 spirals. That analysis was hampered by the inferior angular resolution of the pre-{\\tt COSTAR} HST; furthermore, the Phillips et al.\\ sample was strongly biased toward late-type galaxies. The exploration of the innermost regions of early-type spirals, i.e. of spheroidals at the low-end side of the luminosity sequence, promises to uncover important aspects regarding the formation of nuclei, and their dependence on environment and dynamical structure.  More generally, it may provide a crucial test for bulge -- and galaxy -- formation scenarios.  This paper is organized as follows. In Section 2 we briefly review the criteria that we adopted to choose the sample of galaxies, and the steps of data reduction performed to derive the surface brightness profiles. In Section 3 we present the morphological properties of the 40 new galaxies, the analytical fits used to derive the luminosities and sizes of their bulges, and the procedures adopted to quantify the luminosities and sizes of their nuclear compact sources. In contrast to paper I, in which the bulge-like components were fitted exclusively with a $R^{1/4}$-law, here we fit the inner bulge-like structures with either a $R^{1/4}$-law or an exponential profile, and report the parameters for the best fitting analytical description. In order to ensure homogeneity in the quantification of the bulge properties, we re-performed the analytical fits also for the galaxies presented in paper I (see Appendix A).  In Section 4 we use the enlarged sample of 75 objects to discuss the nuclear versus global properties in early-type spiral galaxies. We summarize our conclusions in Section 5. ",
        "conclusions": "We have analyzed HST {\\tt WFPC2} F606W images for 40 spiral galaxies. The targets belong to the sample of 75 objects introduced in Carollo et al. (1997a). We have studied the optical morphological properties of the new 40 galaxies, derived the surface brightness profiles for 25 of them, and fitted the inner bulge components with an $R^{1/4}$ or exponential law.  By using the enlarged sample of 75 galaxies, we have better quantified the main results presented in paper I, namely: {\\it I.} In about 70\\% of the galaxies there is morphological evidence for the presence of a central component, (morphologically) distinct from the surrounding disk. These ``bulges'' may have an irregular morphology, and occasionally host nuclear star formation.  {\\it II.}  Resolved, central compact sources are detected in about 50\\% of the galaxies. {\\it III.}  The central compact sources in (S0a-Sb) galaxies with nuclear star formation are brighter than those hosted by galaxies with a quiescent nuclear appearance. The luminosity of the compact sources positively correlates with the total galactic luminosity at the 99.9\\% level.  Furthermore, the analysis of the enlarged sample of 75 objects shows that: {\\it I.}  Several of the inner, morphologically-distinct structures are well fitted by an exponential profile.  The subset of these ``exponential bulges'' which are embedded in otherwise normal spiral structure appear to be fainter than $R^{1/4}$ bulges, for constant total galactic luminosity and Hubble type. {\\it II.} The central sources are present in several systems as early as S0a-Sa; about 60\\% of the Sb to Sc galaxies host a central compact source in their very center.  {\\it III.}  The central star clusters are detected in exponential/irregular bulges, or in bulge-less galaxies. Compact sources embedded in $R^{1/4}$-law bulges would possibly not be detected; however, it seems likely that if born, they would be tidally destroyed by the high nuclear densities of these systems, see Carollo \\& Stiavelli 1998).  {\\it IV.}  A large fraction ($\\approx 30\\%$) of the galaxies which host a central compact source shows a barred structure in the optical images.  In several other hosts of compact sources ($\\approx 30\\%$) the presence of a bar is hinted by e.g., an inner square-like morphology. {\\it V.}  Environment seems not to play any role in shaping the nuclear properties of disk galaxies: isolated and not isolated galaxies show the same, and large, variety of properties.  Our main conclusion is that in a significant fraction of galaxies classified from the ground as relatively early-type spirals, there is little or no morphological/photometrical evidence for a classical, smooth, $R^{1/4}$-law bulge. Our results are fully complementary to, and support, studies that have revealed cold kinematics in dense, $R^{1/4}$-law structures embedded in the Sb and later-type spirals (Kormendy 1993): both weaken the connection between ``bulges'' and low luminosity ellipticals, and unveil a large complexity in the inner structure of disk galaxies."
    },
    "9804/astro-ph9804124_arXiv.txt": {
        "abstract": "A one-dimensional turbulent convection model in the form of a time-dependent diffusion equation for the turbulent energy is incorporated into our numerical pulsation code. The effect of turbulent convection on the structural rearrangement of the static equilibrium star is taken into account, and the linear eigenvalues (periods and growth-rates) for the {\\sl complete turbulent convective nonadiabatic} pulsations of the equilibrium models are calculated. The linearized code allows us to perform efficiently a systematic survey of the stability of Cepheid models for a wide range of astrophysical parameters ($L$,$M$,\\Teff) and the turbulent convective parameters.   A sensitivity analysis of the properties of several sequences of Galactic Cepheids is performed with the goal of using observational constraints, such as the shape and location of the fundamental and first overtone instability strips, to calibrate the turbulent convective input parameters.  The locations of the red and blue edges of the fundamental and first overtone instability strips and the overall strip widths are largely determined by the parameters that determine the mixing length, and strengths of the eddy viscosity and of the convective flux.  The remaining parameters can be used for fine-tuning agreement with observations. ",
        "introduction": "Cepheid studies have acknowledged, right from the beginning, that convection should occur in Cepheid envelopes, and that they should alter the properties of both the pulsations and the underlying equilibrium models.  Since the Prandtl number characteristic of stellar material is very small, $\\Pr \\sim O(10^{-9})$, and the Rayleigh numbers associated with the partial ionization regions (PIR) are enormous, $\\Ra$ $\\approxgt$ $O(10^{12})$, stellar convection is far into the hard turbulent regime (Krishnamurti \\& Howard \\cite{Ruby}).  The cooler Cepheid stars are believed to have the more extensive convection zones, leading to pulsational stability beyond the red edge of the standard instability strip (IS).  Because of the estimated inefficiency of convection and because of the difficulties of treating turbulent convection, a lot of work has disregarded it altogether, the hope having been that, except for providing a red edge, convection would play a subdued role.  And purely radiative models have indeed provided a reasonable overall agreement with observations.  More recently it has become increasingly clear that an impasse has been reached with radiative models, and that it is not possible to improve the agreement with observations (e.g. Buchler \\cite{Buchler97}). In particular, the amplitudes of pulsation are systematically too large. No consistent set of pseudo-viscous parameters can correctly limit the amplitudes of {\\sl both} fundamental and first overtone pulsations without adversely affecting the stability of the corresponding limit cycles, and the agreement with observation.  The new EROS and MACHO data of the low metallicity Magellanic Cloud Cepheids (Buchler et al. \\cite{BKBG}) reveals further disagreement between stellar evolution and radiative stellar pulsation models.  Clearly some additional physical dissipation is missing in the radiative codes (see also Kov\\'acs \\cite{Kovacs}), and turbulence and convection are the primary suspects.  Simulation of even the simplest turbulent convection (TC) problem is a formidable numerical challenge; to adequately resolve the many time- and length-scales of fully turbulent 3D convection at stellar Rayleigh and Prandtl numbers is totally out of question, and astrophysicists have been looking for 1D recipes to give an approximate, but acceptable treatment of this phenomenon.  Thus the 1925 mixing length picture of Prandtl was readily adapted to stellar envelopes by B\\\"ohm-Vitense, and subsequently reformulated in many variations, all of which are equivalent for time-independent problems (for an overview, see Baker \\cite{Baker}).  But time-independent MLT was never able to resolve the red edge problem, mostly because it is time-dependent dissipation introduced by eddy viscosity that provides a clear red edge, as we will show in \\S~\\ref{limsection}. Spiegel \\cite{Ed63} first extended MLT to the time-dependent and nonlocal cases.  The earliest attempts to include a time-dependent mixing length model in pulsation (Gough \\cite{Gough}, Unno \\cite{Unno}) were not successful because they were too local in space.  Subsequently, Gough's theory was extended to be less local, and eventually was used in a full linearization of Solar modes (Balmforth \\cite{Neil}), analogous to what we have done here.  Unno's theory was also developed (Castor \\cite{Castor}) as a diffusion model, and simplified to a single equation for the turbulent energy $e_t$ by Stellingwerf (\\cite{Swf82}).  More recent applications of this and similar turbulent diffusion recipes have been made by Gonczi \\& Osaki \\cite{Gonc-O}, Bono \\& Stellingwerf \\cite{BS}, Gehmeyr \\& Winkler \\cite{Geh-W}, hereafter GW) and Kuhfuss (\\cite{KuhfuB}).  Further closely related models also appear in the literature, such as those of Canuto (\\cite{Canuto}) but have yet to be applied.  All 1D model TC equations contain several dimensionless, order unity parameters that are directly related to the physical quantities of the model.  In principle, the validity of the 1D model equation can be checked against experiments or against detailed 3D simulations of turbulent convection, and values of the unknown parameters can be extracted.  However, neither such experiments nor such simulations are presently available.  We are obliged to calibrate these parameters indirectly with the help of observational astronomical constraints, as suggested more than a decade ago by Stellingwerf (\\cite{Swf82}).  In the best outcome, observational constraints could select one model in favor of the others.  On the other hand we might discover that single equation models for turbulent convection are insufficient, or that plumes (e.g. Rieutord \\& Zahn \\cite{R-Zahn}) play an essential role in the convective transport.  This programme has recently been initiated by incorporating these recipes into hydrodynamic pulsation codes.  Fully nonlinear models have been used to determine the stability properties of convective pulsators (e.g. Bono \\& Stellingwerf \\cite{BS}, Bono \\& Marconi \\cite{BM97}), but this approach is inefficient, and unsuitable when a large number of models is required.  The approach we follow in this paper is to compute the linear nonadiabatic properties directly by linearizing the hydrodynamic equations about the equilibrium model and by then computing the linear eigenvalues, a procedure which is relatively fast. This allows us to perform an extensive survey of the sensitivity to the TC parameters. ",
        "conclusions": "Purely radiative Cepheid models cannot satisfactorily account for all observational data, and this has forced us to include turbulence and convection in our Cepheid models.  We have adopted a simple 1D model diffusion equation for turbulent convection (TC) that can easily be incorporated into a radial stellar pulsation code.  Although there is no guarantee that such a single 1D equation can satisfactorily approximate the effects of turbulence and convection it is worth exploring to what extent we can obtain agreement with observation. The model TC equations contain several dimensionless, order unity parameters that are directly related to the physical quantities of the model that need to be calibrated with the help of observational astronomical constraints.  Instead of performing time-consuming fully nonlinear hydrodynamic computations of the pulsations, we have developed a code to perform an efficient linear nonadiabatic stability analysis of the models.  This has allowed us to make an exploration of the sensitivity to the free parameters of the TC equation, both of the structure of the equilibrium Cepheid models and of their linear properties.  In agreement with other studies we find that the static Cepheid models exhibit convection primarily in their H and He partial ionization regions (PIRs), although in some parameter range a convective zone also appears in the Fe PIR.  For strongly convective models the separate convective zones merge into a large one that penetrates up to 320,000\\th K.  The coupling of pulsation and convection also modifies the linear properties of the pulsational modes, as expected.  It also creates a new branch of turbulent diffusion modes linear that our stability analysis establishes to be extremely damped, fortunately for our 1D TC recipe. These additional modes thus cannot create havoc numerically or physically. This large damping is consistent with the convective timescales being very short compared to the periods of the vibrational of the excited modes which themselves are short compared to the growth-times.  Our survey shows that the stability of the fundamental and first overtone modes depends on all the TC parameters, but that it is dominated by (1) the convective flux, (2) the eddy viscosity, and (3) the mixing length.  Generally, the effects of eddy viscosity and mixing length are greater in the higher overtones.  In agreement with other much less detailed TC studies we find that reasonable $\\alpha$ parameter values exist for which IS's of the galactic Cepheid F and O1 have an acceptable range of temperatures.  In this paper we have made no effort to narrow down the acceptable range of TC parameters by using the large amount of observational information. An application of the TC code to a broader sample of Cepheid models and its comparison to the large data base of Galactic, and Magellanic Cloud Cepheids is in progress, as is an extension to RR~Lyrae stars."
    },
    "9804/astro-ph9804254_arXiv.txt": {
        "abstract": "\\noindent  We present the results of an exhaustive numerical study of fully relativistic non-axisymmetric Bondi-Hoyle accretion onto a moving Schwarzschild black hole. We have solved the equations of general relativistic hydrodynamics with a high-resolution shock-capturing numerical scheme based on a linearized Riemann solver. The numerical code was previously used to study axisymmetric flow configurations past a Schwarzschild hole. We have analyzed and discussed the flow morphology for a sample of asymptotically high Mach number models. The results of this work reveal that initially asymptotic uniform flows always accrete onto the hole in a stationary way which closely resembles the previous axisymmetric patterns. This is in contrast with some Newtonian numerical studies where violent flip-flop instabilities were found. As discussed in the text, the reason can be found in the initial conditions used in the relativistic regime, as they can not exactly duplicate the previous Newtonian setups where the instability appeared. The dependence of the final solution with the inner boundary condition as well as with the grid resolution has also been studied. Finally, we have computed the accretion rates of mass and linear and angular momentum. ",
        "introduction": "In a previous paper (Font and Ib\\'a\\~nez 1997, hereafter Paper I) we studied the morphology and dynamics of relativistic Bondi-Hoyle accretion in axisymmetric flows past a Schwarzschild black hole. The main conclusion of that work was to extend the validity of the Bondi-Hoyle accretion picture (Hoyle and Lyttleton, 1939; Bondi and Hoyle, 1944) to the relativistic regime, finding that the matter is always accreted onto the hole in a stationary way. Furthermore, if the flow was, initially, supersonic, the main feature of the accretion pattern was the presence of a shock cone in the solution. At the same time, we checked the validity of our numerical code revisiting an existing previous calculation (Petrich et al., 1989) using more accurate numerical techniques specifically designed to capture discontinuities. In the present investigation we have extended those studies to account for non-axisymmetric configurations. We have only considered uniform flows at infinity focussing on studying whether or not the flow pattern reaches, ultimately, a stationary state. The main purpose of this work is, then, to extend previous non-axisymmetric Newtonian computations to the realm of general relativity. As far as we know, this has not been studied so far. In particular, we plan to address if non-steady patterns, so prominent in some of the 2D non-axisymmetric simulations in the classical regime, also arise here.  During the last years, a large number of wind accretion simulations past a gravitating body have been performed in Newtonian hydrodynamics (see, e.g., the list of references in Paper I). One of the most interesting features, only revealed by numerical simulations, was the appearance of unstable accretion patterns, contrary to the theoretical (and simplified) Bondi-Hoyle (1944) accretion picture. This highly non-steady behaviour on the wake of the accretor was only found in 2D non-axisymmetric simulations, especially when the resolution employed was fine enough. These unstable patterns are characterized by the shock wave moving from side to side (the so-called flip-flop instability) and also by the appearance of transient phases of disc formation where the angular momentum accreted by the central object significantly increases. This kind of behaviour has been found not only assuming local density and velocity gradients (Fryxell and Taam, 1988; Taam and Fryxell, 1989) but also when considering accretion of uniform flows at infinity (Matsuda et al., 1991). However, in detailed 3D computations performed more recently (Ruffert and Arnett, 1994; Ruffert, 1994a,b, 1995, 1996) the accretion cone remains quite stable and no sign of flip-flop instability appears unless the flow is assumed to have explicit density gradients at infinity (Ruffert, 1997).  In all previous studies it was shown that the key parameter which controls the appearance of the instability was the {\\it size of the accretor}. In particular, this size should be a very small fraction of the accretion radius in order to find any evidence of flip-flop. For instance, in Sawada et al. (1989) and Matsuda et al. (1991) computations, the flip-flop instability appeared, in 2D, for a central size object of $\\approx 0.0625 r_a$, where $r_a$ is the accretion radius, the natural length scale of this problem, defined as  \\begin{equation} r_{a} = {{2GM}\\over {v_{\\infty}^{2}}}. \\end{equation}  \\noindent Here, $M$ is the mass of the accreting object, $G$ is the gravitational constant and $v_{\\infty}$ is the asymptotic velocity of the fluid. They also found that the flow was stable if the size was larger than $\\approx 0.125 r_a$. More recently, Benensohn, Lamb and Taam (1997) have performed detailed 2D non-axisymmetric computations with a smaller accretor of size $0.0375 r_a$ finding unstable behaviour as well.  In addition, there are other parameters that, combined with the size of the accretor, can play a role in the development of the instability. In particular, one of these is the {\\it Mach number} of the flow. The larger the Mach number at infinity is, the more turbulent the wake becomes. Finally, another extremely important consideration is the {\\it position of the shock}, namely if it is attached (tail shock) or detached (bow shock) with respect to the central object. This is basically controlled by the values of the asymptotic Mach number and the adiabatic index of the gas. For low Mach number and high adiabatic index flows (say $5/3$), the shock is mainly detached at the front while the opposite trend produces a tail shock. All previous studies where the instability appeared needed the formation of a detached shock in front of the accretor. As the evolution procceeded, this detached shock transformed into a dome-shaped shock and, eventually, this was followed by a violent flip-flop instability (Sawada et al., 1989; Matsuda et al., 1991; Livio et al., 1991; see also Ruffert, 1997). In particular, Livio et al. (1991) demonstrated, using a simplified analytic treatment, the existence of a wake instability to small deflections when the shock opening angle is large. This can happen even for an initially homogeneous flow if the accretion process has low efficiency. However, the flow can be stabilized if the accreting object is larger than some critical dimension (Livio et al., 1991).  According to these previous classical results, there are several factors that can lead to the non-appearance of this instability in the relativistic regime when the accretor is a black hole. First, and most important, in this regime there is a physical minimum value for the size of the accreting object, given by the gravitational (Schwarzschild) radius. This, of course, places some constraints in the possibility of extending previous Newtonian studies to very small values of the accretor, assuming that we also need an asymptotically supersonic flow (in order to have any shock wave) and avoiding, at the same time, flow velocities larger than the speed of light. The second factor has to do with the position of the shock. As we will show below (see also Paper I), all models that we have evolved are characterized by the presence of a {\\it tail} shock in the rear part of the hole. Hence, according to the Newtonian studies, this would not help the development of the flip-flop instability.  The paper is organized as follows: in next Section ($\\S 2$) we briefly present the system of equations of general relativistic hydrodynamics written as a hyperbolic system of conservation laws. The numerical code and the initial setup are also described here. The results of the simulations are presented and analyzed in Section $\\S 3$. Finally, Section $\\S 4$ summarizes the main conclusions of this work. ",
        "conclusions": "We have performed a detailed numerical study of relativistic Bondi-Hoyle accretion onto a moving Schwarzshild black hole. We have extended our previous work to include non-axisymmetric configurations. This has been motivated by the possibility of finding unstable accretion patterns in the flow. We have only considered flows which are initially uniform. We have found that, contrary to Newtonian simulations, the non-axisymmetric solution is always stable to tangetial oscillations. However, this statement should be relaxed, as the corresponding initial setup in our relativistic wind simulations can not reproduce, for physical reasons, the one used in classical simulations. In particular, we have considered values of the accretion radius which, when used in the classical regime, would have not given rise to any instability as well. This may indicate that, in the relativistic regime, a uniform flow accreting onto a black hole could always remain symmetric and stationary. It might be worthwhile to go to much larger values of the asymptotic Mach number to see if any unstable pattern actually develops. These are, however, quite challenging simulations that would push the code to its limits, as the asymptotic sound speed should be considerably small.  We have checked the stability and stationarity of our relativistic accretion patterns by analyzing the broad morphology of the flow with isocontour plots as well as by integrating, over all time of the simulation, global quantities as the accretion rates of mass and momentum. In all cases considered, we have found that the mass and radial momentum accretion rates approach asymptotically to very accurate constant values in time. In addition, the angular momentum accretion rate always remains in very low values, indicating the total absence of rotating flows around the hole which could have triggered any tangetial instability.  It maybe worthwhile to mention that, though we have restricted in the present work to homogeneous flows, the code is capable of evolving more complicated non-axisymmetric configurations where the flow has some amount of angular momentum at infinity. We are currently working in a parametric study of this kind of flows (Pons et al., 1998). In particular, we plan to gradually increase the angular momentum of the incident flow and study how it is transfered onto the hole via alternate spiral shocks. At the same time, we also plan to add rotation to the hole -- a feature that the actual code used in this work already incorporates -- extending thus the relativistic Bondi-Hoyle accretion to moving Kerr black holes. Ultimately, we plan to perform these simulations in three spatial dimensions."
    },
    "9804/astro-ph9804018_arXiv.txt": {
        "abstract": "The Hubble Deep Field South (HDF-S) Hubble Space Telescope (HST) observations are expected to begin in October 1998.  We present a composite spectrum of the QSO in the HDF-S field covering UV/optical/near IR wavelengths, obtained by combining data from the Australian National University 2.3m Telescope with STIS on the Hubble Space Telescope (HST)\\footnote{Our ground-based 2.3m spectrum and the composite spectrum are available in electronic format at the address http://bat.phys.unsw.edu.au/$\\sim$kms/hdfs/}.  This intermediate resolution spectrum covers the range 1600--10000\\AA~ and allows us to derive some basic information on the intervening absorption systems which will be important in planning future higher resolution studies of this QSO.  The QSO J2233$-$606 coordinates are 22:33:37.6 $-$60:33:29 (J2000), the magnitude is B=17.5, and its redshift is z$_{em} = 2.238$, derived by simultaneously fitting several emission lines.  The spectral index is $\\alpha = -0.7 \\pm 0.1$, measured between the Lyman-$\\alpha$ and Mg II emission lines.  Many absorption systems are present, including systems with metal lines redwards of the Lyman-$\\alpha$ emission line at $z_{abs} = 2.204, 1.942, 1.870, 1.787$ and a few very strong Lyman-$\\alpha$ features at $z_{abs} = 2.077, 1.928$, without similarly strong metal lines.  There is a conspicuous Lyman limit (LL) absorption system that is most likely associated with the $z_{abs} = 1.942$ system with a neutral hydrogen column density of N$_{HI}=(3.1\\pm1.0)\\times10^{17} cm^{-2}$.  There is some evidence for the presence of a second LL absorber just to blue of the conspicuous system at $z = 1.870$.  We have employed a new technique, based on an analysis of the shape of the observed spectrum in the region of the LL absorption, to explore the properties of the gas.  We tentatively conclude that this system might have suitable characteristics for measuring the deuterium-to-hydrogen (D/H) ratio. ",
        "introduction": "Compared to the original HDF, scientific studies of the the Hubble Deep Field South (HDF-S) may benefit from the additional information provided by a bright high redshift $z>2$ QSO lying close to (but not within) the targeted WFPC-2 imaging area.  Initially 12 small ($1\\deg~\\times~1\\deg$) fields were selected as candidate target fields.  Whilst no known high $z$ QSOs were present in any of the possible fields, Hewett and Irwin found a high $z$ QSO candidate in one of the fields using a UK Schmidt Telescope IIIaJ objective prism plate scanned by the Automated Plate Measuring facility in Cambridge. Follow-up observations of the candidate at the Anglo-Australian Telescope confirmed the object to be a QSO (J2233$-$606 22:33:37.6 $-$60:33:29 (J2000), B=17.5) with a redshift $z>2$ (see Boyle 1997 for details).  We note for completeness that 20cm radio observations have been made (Ray Norris, private communication reported in Boyle 1997) that give an upper limit on the radio flux as S$_{20cm}<3$mJy.  There is a strong radio source 10 arcmin from the QSO (S$_{20cm}=146$mJy).  We present here a low-resolution spectrum of the HDF-S QSO taken with the Australian National University (ANU) 2.3m Telescope and combined with the publicly available STIS spectrum providing wide spectral coverage.  The QSO has an optically thin Lyman limit which is a possible candidate for measuring the deuterium-to-hydrogen (D/H) ratio. The QSO also has several strong absorption features redwards of the Lyman-$\\alpha$ emission line, to which we assign tentative identifications. ",
        "conclusions": "We have reported the results of a study of a low/intermediate resolution spectrum covering the UV/optical/near IR of the HDF-S QSO J2233$-$606, formed by combining HST/STIS observations with higher wavelength data from the ANU 2.3m telescope.  The intrinsic QSO spectral characteristics have been measured from this data.  A study of the absorption properties have revealed multiple CIV absorption systems, suggesting that either a large structure or multiple galaxies or clusters intersect the sightline to this QSO. Structure in the LL absorption is probably caused by the presence of at least 2 high N(HI) systems, one of which may have properties which make it suitable for studying D/H.  Despite the low spectral resolution, we have been able to derive reasonably detailed characteristics of the absorption systems.  Future high resolution spectra will of course yield far more information. However, it is already clear from our results that this QSO and the HDF-S studies of its environs will reveal a host of fascinating new information about the z$\\sim 2$ Universe.  {"
    },
    "9804/astro-ph9804197_arXiv.txt": {
        "abstract": "\\object{26W20}, a radio galaxy located at a projected distance of 2.2 Mpc from the center of Abell 754 (z=0.054) exhibits core emission properties similar to those of a BL Lac object.  New radio observations and VLA archival data show parsec scale features in the core (\\(L_{r}\\sim10^{41}\\) ergs s$^{-1}$) with a flat spectrum.  From observations with ROSAT, we demonstrate that the X-ray core emission \\((L_{x}=3\\times10^{43}\\) ergs s$^{-1}$) is non-thermal and variable.  An 18\\% increase in the X-ray luminosity was observed in two observations separated by five days. The optical morphology is that of a large, normal elliptical galaxy (\\(m_{v}=15.6)\\) and the optical spectrum lacks strong emission lines.  If the innermost segment of the radio jet is close to our line of sight, Doppler boosting could explain both the observed X-ray intensity and the absence of emitting line regions. ",
        "introduction": "26W20 was discovered in the Westerbork radio survey by Harris {\\it et al.} (1980).  The galaxy lies in a sub-cluster of \\object{Abell 754}.  The structure of the cluster and the reality of the sub-cluster containing 26W20 are detailed by Kriesslar and Beers (1997).  The radio morphology and classification has been described as a Fanaroff and Riley class I object, either a 'tailed radio galaxy' (TRG) or a single-sided lobe (Harris {\\it et al.} 1980; Harris, Costain and Dewdney, 1984 - 'HCD' hereafter).  A single jet/tail is seen extending 170 kpc from the core (\\(H_{o}\\)= 50 km s$^{-1}$ Mpc$^{-1}$) and ending in a low surface brightness lobe (Fig. 1). Because no jet or lobe was detected on the other side of the galaxy, and since the host galaxy is a member of a small group of galaxies, a TRG model was favored over a single-sided jet model (HCD).  Galaxies with X-ray bright nuclei are usually characterized by strong emission lines.  26W20 is one of four known galaxies with a high apparent X-ray luminosity from the nucleus but with only weak or undetectable emission lines.  The other three are \\object{3C264} (Baum {\\it et al.}  1997), \\object{J2310$-$43} (Tananbaum {\\it et al.} 1997) and \\object{PKS 2316$-$423} (Crawford and Fabian 1994).  Nuclear parameters are given in Table 1.  26W20 was first observed in X-rays with the EINSTEIN Observatory's imaging proportional counter (IPC) and high resolution imager (HRI). The X-ray source was unresolved with the HRI (FWHM=5\\arcsec) and the IPC spectrum could be equally well fit by a power law with an energy index of 0.8 \\(\\pm\\)0.4 or a thermal bremsstrahlung spectrum with kT=3 keV (constrained to be greater than 1 keV).  Motivated to understand the head-tail morphology of 26W20, we have obtained new radio and X-ray measurements to study the inner jet region.  While there is possible evidence for a TRG model, these observations reveal core emission properties similar to those of a BL Lac object. However, there is no evidence for excess optical emission from the nucleus.  This distinguishes 26w20 from a typical BL Lac object.  We use ROSAT HRI data to provide upper limits on hot ambient gas which would be required to provide the 'ram pressure' if nuclear gas is to be 'stripped' from the galaxy. MERLIN observations were obtained to see if twin jets are present in the core and to see if they show severe bending necessary for a beaming model in order to change the jet from being close to the line of sight to that observed for the large scale tail which lies more in the plane of the sky.  \\scriptsize \\begin{flushleft} \\begin{tabular*}{90mm}{lllll} \\multicolumn{5}{c}{Table 1.} \\\\ \\multicolumn{5}{c}{Comparative parameters for the nucleus of similar sources} \\\\ \\hline\\\\ Parameter&26W20&J2310-43$^{1}$&3C 264$^{2}$&PKS 2316-423$^{3}$\\\\ \\hline Redshift..............&0.054&0.0886&0.0215&0.0549\\\\ l$_{r}$(10$^{30}$ cgs)$^{a}$&7.8&21.2&5.2&70.3\\\\ L$_{x}$($10^{43}$ergs s$^{-1}$).&3.7$^{b}$&14.5$^{c}$&0.88$^{c}$&8.6$^{c}$\\\\ M$_{V}^{d}$.....................&$>$-19.6&$>$-20.1&$>$-19.8&$>$-20.8$^{e}$\\\\ $\\alpha_{r}$......................&0.25&$\\sim$0&$\\sim$0&0.67\\\\ $\\alpha_{x}$......................&1.3&1.4&1.3&---\\\\ Emission Lines...&none&none&weak$^{f}$&weak\\\\ \\label{sources} \\end{tabular*} \\end{flushleft} References: (1) as reported in Tananbaum et al. 1997 (2) as reported in Tananbaum et al. 1997, Elvis et al. 1981; (3) Crawford and Fabian 1994\\\\ \\\\ Notes: (a) Monochromatic luminosity at 5 GHz in units of ergs s$^{-1}$ Hz$^{-1}$; (b) Energy range of (0.5-3.0 keV); (c) Energy range of (0.1-2.4 keV); (d) The core magnitude is assumed to be less than 1/10 of the magnitude of the whole galaxy; (e) This lower limit is due to the inclusion of optical emission not associated with the nucleus; (f) An extended region of 400 km s$^{-1}$ wide H$\\alpha$+[NII] emission is seen (Baum {\\it et al.}  1997); \\normalsize ",
        "conclusions": "The radio and X-ray core emission in conjunction with the lack of prominent optical emission lines are properties similar to a BL Lac type object but the lack of strong optical core emission distinguishes this galaxy from a typical BL Lac.  Our high resolution MERLIN observation at 6 cm (Figure \\ref{6cm}) shows a compact region of about 100 parsecs. These results further support the compact nature of the core emission of the VLA images by Harris {\\it et al.} (1984) and the archival 3.6 cm data (Figure \\ref{4cm}).  The position angle of the second component (PA=140\\degr) agrees, within the errors, with those in the 21 cm map (PA=145\\degr) and the 3.6 cm image (PA=150\\degr).  We have computed the values for $\\alpha_{ro}$(0.43) and $\\alpha_{ox}$(0.99) with an assumption that the upper limit of the optical core flux is 0.24 mJy (1/10 of the total flux of the galaxy in the V band).  The X-ray flux from an HRI measurement is 0.56 $\\mu$Jy at 1 keV.  Given these values, 26W20 lies in the center of the X-ray selected BL Lac objects on the $\\alpha_{ro}$ vs. $\\alpha_{ox}$ plot as shown in Figure 8 of Tananbaum {\\it et al.} 1997.  If the core luminosity is significantly less than our upper limit, this could place 26W20 at an extreme position outside of the BL Lac group.  Given the range of $\\alpha_{ro}$ and $\\alpha_{ox}$ for a typical X-ray selected BL Lac, the core optical flux is between 3.2$\\times10^{-2}$ and 1.55 mJy.  Therefore, an excess of optical flux in the core will not significantly change the color or magnitude of the galaxy as a whole.  No evidence for the presence of any hot, extended gas was found in the sub-cluster containing 26W20.  From density measurements of the sub-cluster environment, the values fall within the range of ambient densities of other clusters harboring TRGs.  We propose that the BL Lac properties observed in 26W20 are evidence of energetic emission from the core due to beamed emission from a component of the jet.  Therefore, the large apparent (i.e. if emission is isotropic) X-ray luminosity of $10^{43}$ ergs s$^{-1}$, concomitant with the absence of emission lines, means either there is no gas in the galaxy to form emission line regions or the X-rays are beamed and not available to ionize the gas."
    },
    "9804/astro-ph9804183_arXiv.txt": {
        "abstract": "We have calculated \\ion{Fe}{2} emission line strengths for Active Galactic Nuclei Broad-Line Regions using precise radiative transfer and Iron Project atomic data. We improve the treatment of all previously considered excitation mechanisms for the \\ion{Fe}{2} emission, continuum fluorescence, collisional excitation, fluorescence by self-overlap among the iron lines, and fluorescent excitation by \\lya. We demonstrate that \\lya\\ fluorescence is of fundamental importance in determining the strength of the \\ion{Fe}{2} emission. In addition to enhancing the ultraviolet and optical \\ion{Fe}{2} flux, \\lya\\ fluorescence also results in significant near-infrared \\ion{Fe}{2} emission in the $\\lambda\\lambda8500-9500\\,$\\AA\\ wavelength range. New observations are suggested to probe this effect in strong \\ion{Fe}{2} emitting quasars. ",
        "introduction": "One of the most puzzling instances of \\ion{Fe}{2} line formation occurs in the spectra of Active Galactic Nuclei (AGN) with Broad-Line Regions (BLR) which exhibit strong ultraviolet \\ion{Fe}{2} emission (UV$\\,\\equiv \\lambda\\lambda2000-3000\\,$\\AA) and often strong optical \\ion{Fe}{2} emission (OP$\\,\\equiv \\lambda\\lambda3000-6000\\,$\\AA) as well. The widths of the iron lines indicate formation in the dense BLR, yet standard photoionization calculations, successful in interpreting many other BLR features, fail to account for the strength of the \\ion{Fe}{2} emission (Wills, Netzer \\& Wills 1985; Joly 1993). Observed total \\ion{Fe}{2}/H$\\beta$ ratios are $\\approx 12$, whereas standard models predict ratios less than 8, typically 3 - 5, with the largest values resulting from enhanced iron abundances. There is also a class of ``superstrong'' \\ion{Fe}{2} emitters with \\ion{Fe}{2}/H$\\beta \\approx 30$ which seem unable to be explained by traditional photoionization modeling (Joly 1993; Graham \\etal\\ 1996).  Excitation mechanisms advanced for the \\ion{Fe}{2} emission include continuum fluorescence via the UV resonance lines, self-fluor\\-es\\-cence via overlapping \\ion{Fe}{2} transitions, and collisional excitation, with the latter thought to contribute the bulk of the \\ion{Fe}{2} emission. However, Wills \\etal\\ (1985) demonstrated that model BLR calculations including these excitation mechanisms cannot account for the observed strengths of the \\ion{Fe}{2} lines. Elitzur \\& Netzer (1985, hereafter EN) extended this calculation to include \\lya\\ fluorescent excitation but found its contribution to be negligible. Nevertheless, Penston (1987) argued that strong indirect observational evidence for the importance of \\lya\\ fluorescence is found in the presence of unexpected UV \\ion{Fe}{2} multiplets in the spectrum of the symbiotic star RR Tel which seem attributable only to cascades from higher levels pumped by \\lya\\ fluorescence. The emission nebulae of symbiotics offer densities and ionization parameters similar to the BLRs of AGN. Recently, Graham \\etal\\ (1996) have identified emission from these UV \\ion{Fe}{2} multiplets expected to be preferentially strengthened by \\lya\\ in the spectrum of the ultra-strong \\ion{Fe}{2} emitter 2226-3905.  In this {\\it Letter}, we improve the modeling of all of the mentioned excitation mechanisms by using new, accurate atomic data and precise radiative transfer, and we demonstrate that fluorescent excitation by \\lya\\ can be of prime importance. A key feature of \\lya\\ fluorescent excitation of \\ion{Fe}{2} is significant near-IR \\ion{Fe}{2} emission in the wavelength range $\\lambda\\lambda8500-9500\\,$\\AA. We suggest that a currently well-known, but unidentified, emission feature near $\\lambda9200$ in AGN spectra may be, at least in part, a spectral signature of this process. ",
        "conclusions": "Using precise radiative transfer and accurate atomic data for \\ion{Fe}{2}, we demonstrate that \\lya\\ fluorescent excitation can more than double the \\ion{Fe}{2} flux in the UV and optical regions. We also predict \\ion{Fe}{2} emission in the $\\lambda\\lambda8500-9500\\,$\\AA\\ region as a direct result of \\lya\\ fluorescence. Both of these results can be viewed as theoretical confirmation of the role of \\lya\\ fluorescence suggested by Penston (1987) on the basis of observations. New observations of the $\\lambda\\lambda8500-8500\\,$\\AA\\ wavelength range in strong \\ion{Fe}{2} emitters are highly desirable, and would provide valuable constraints on modeling $\\lya$ fluorescence. Potentially new information on the \\lya\\ profile within single BLR clouds could be obtained, something which is impossible based on observations of the highly broadened \\lya\\ profile itself.  A detailed discussion of the variation of the \\ion{Fe}{2} emission with the physical parameters defining the BLR clouds (using a much larger 827 level \\ion{Fe}{2} atom which includes the doublet spin system), and the implications for the `\\ion{Fe}{2}/H$\\beta$' problem, will be presented by Sigut \\& Pradhan (1998).  We are also extending our calculations to overcome some of the limitations of the present work, in particular, to incorporate all important ions into a full solution of the non-LTE radiative transfer problem for the thermal structure of AGN BLR clouds. This will enable the \\lya-\\ion{Fe}{2} problem to be solved within a self-consistent framework."
    },
    "9804/astro-ph9804076_arXiv.txt": {
        "abstract": "The current census of observed baryons in the local Universe is still missing a significant fraction of them according to standard Big-Bang nucleosynthesis. Numerical simulations predict that most of the missing baryons are in a hot intergalactic medium, which is difficult to observe through its X-ray emission or Sunyaev-Zel'dovich effect.  We show that the next generation of X-ray satellites will be able to detect this gas through the X-ray absorption lines imprinted by its highly-ionized metals on the spectrum of a background quasar. For the metallicity typically found in intracluster gas, up to 70\\% of the baryons produce O VIII absorption lines with an equivalent width $\\ga 0.1$ eV. The spectrum of any high redshift quasar is expected to show several such lines per unit redshift due to intervening gaseous halos of galaxy groups. These lines will be detectable at a signal-to-noise ratio $\\ga 5$ after a day of integration with the future Constellation-X telescope for any of the several hundred brightest quasars across the sky. ",
        "introduction": "The best estimate for the sum of the baryonic mass currently observed in stars, neutral hydrogen, and X-ray emitting gas in clusters of galaxies (Fukugita, Hogan, \\& Peebles 1998) falls short of the most-likely value predicted by standard Big-Bang nucleosynthesis (Burles \\& Tytler 1998). Numerical simulations suggest that most of the missing baryons reside in a hot intergalactic medium, which is too rarefied to be detectable through its X-ray emission or the Sunyaev-Zel'dovich effect (e.g., Ostriker \\& Cen 1996; Theuns et al.  1998).  Much of this gas resides in the outskirts of galaxy groups or clusters, and hence is likely to possess a similar metallicity to that observed in the X-ray emission spectrum from the cores of these systems (Mushotzky et al.  1996).  In this paper we explore the possibility of detecting this hot intergalactic medium through the X-ray absorption lines imprinted by its metals on the spectrum of a background quasar.  Detection of X-ray absorption lines from galaxy clusters had been suggested in the past as a means of measuring distances (Krolik \\& Raymond 1988; Sarazin 1989).  However, because of the limited sensitivity of X-ray instruments at that time, these early papers restricted their attention to the strong lines produced in the cores of rich clusters and hence to the unlikely superposition of a quasar behind such cores. The next generation of X-ray telescopes, manifested at their best by the successor to AXAF, Constellation-X (see http:// constellation.gsfc.nasa.gov), will achieve a much greater sensitivity than previously imagined. It is therefore the purpose of this paper to quantify the full probability distribution of equivalent widths for the intergalactic X-ray absorption lines imprinted on the spectrum of any {\\it random} quasar.  Since much of the baryonic mass is currently assembled into virialized objects and since the present nonlinear mass-scale is that of galaxy groups, most of the absorption lines would be produced by gas in the outskirts of these groups.  In \\S 2 we combine the Press-Schechter mass function and the singular-isothermal-sphere model to predict the density and temperature distribution of this gas in a standard CDM cosmology with $\\sigma_{8h^{-1}}=0.5$.  We used a Bardeen et al.  (1986) power spectrum with a tilted primordial index of $n=0.7$. Because the power-spectrum is normalized on the cluster mass scale, our results depend only weakly on the choice of the power spectrum shape (and agree by construction with X-ray data for the temperature function of groups and clusters; see, e.g. Henry et al.  1995).  We assume an $\\Omega=1$, $\\Omega_b=0.1$, $h=0.6$ cosmology throughout the discussion.  In \\S 3 we compute the corresponding statistics of the absorption lines produced by highly-ionized metals in this gas. Finally, \\S 4 summarizes the main conclusions of this work. ",
        "conclusions": "We have investigated the possibility of using the X-ray absorption spectrum of a background quasar to probe the presence of heavy elements in the outer regions of groups and clusters of galaxies.  Figure 1 demonstrates that observations of the O VIII line with an equivalent width sensitivity $\\ga 0.1$ eV should be able to probe most of the baryonic material out to the virial radius of these systems, well beyond the regions where X-ray emission is currently detected.  We estimate that there should be several such O VIII absorption lines per unit redshift along the line of sight to any quasar (cf. Fig. 2).  The future Constellation-X telescope will be able to detect these lines with a signal-to-noise ratio $\\ga 5$ in a day of integration for any of the few hundred brightest quasars accross the sky.  Our model assumed that the gas in groups or clusters of galaxies is isothermal. The modest radial decline which was observed recently with ASCA for the temperature of the gas in the outer parts of clusters (Markevitch et al.  1998), might make absorption features by ionic species such as O VII more prominent than O VIII in the X-ray absorption spectra of quasars. On the other hand, energy input by the same supernovae that enrich the gas with metals might increase the temperature of X-ray halos beyond the value adopted in equation~(\\ref{eq:M-T}) (which was calibrated based on numerical simulations of gravitational dynamics only).  In our calculations we have assumed a constant metallicity of the gas; a reduced metal abundance due to inefficient mixing in the outer parts of groups and clusters would obviously lower our predicted signal. If absorption will be detected for some quasars, it would be interesting to correlate the absorption line redshift and sky position with X-ray or optical catalogs of galaxy clusters. If a group or a cluster are identified at the absorption redshift, it would be possible to infer the impact paramter $b$, and the metallicity $Z$, using equations~(\\ref{eq:Ni}) and~(\\ref{eq:W}).  Observations of the O VIII line around X-ray clusters at different redshifts could trace the distribution of metals in the intergalactic medium and possibly shed light on the epoch and nature of the metal enrichment process."
    },
    "9804/astro-ph9804240_arXiv.txt": {
        "abstract": "We have obtained radial gradients in the spectral features \\mbox{$\\lambda$4000-\\AA} break (D$_{4000}$) and Mg$_2$ for a sample of 11 central cluster galaxies (CCGs): 8 in clusters with cooling flows and 3 in clusters without. After careful removal of the emission lines found within the D$_{4000}$ and the Mg$_2$ bandpasses for some objects, the new data strongly confirm the correlations between line-strength indices and the cooling flow phenomenon found in our earlier study (Cardiel, Gorgas \\& Arag\\'{o}n-Salamanca 1995).  We find that such correlations depend on the presence and characteristics of emission lines in the inner regions of the CCGs. The nuclear indices are correlated with the mass deposition rate ($\\dot{M}$) {\\it only\\/} when emission lines are found in the central regions of the galaxies. The central D$_{4000}$ and Mg$_2$ indices in cooling flow galaxies {\\it without\\/} emission lines are completely consistent with the indices measured in CCGs in clusters without cooling flows.  CCGs in cooling flow clusters exhibit a clear sequence in the \\mbox{D$_{4000}$--Mg$_2$} plane, with a neat segregation depending on emission-line types (Heckman et al.  1989) and blue morphology (McNamara 1997). This sequence can be modelled, using stellar population models with a normal IMF, by a recent ($\\sim 0.1$~Gyr old) burst of star formation, although model uncertainties do not allow us to completely discard a continuous star formation or a series of bursts over the last few Gyr.  In CCGs with emission lines, the gradients in the spectral indices are flat or positive inside the emission-line regions, suggesting the presence of young stars.  Outside the emission-line regions, and in cooling flow galaxies without emission lines, gradients are negative and consistent with those measured in CCGs in clusters without cooling flows and giant elliptical galaxies.  Index gradients measured exclusively in the emission-line region correlate with $\\dot{M}$. Using the same population models we have estimated the radial profiles of the mass transformed into new stars.  The derived profiles are remarkably parallel to the expected radial behaviour of the mass deposition rate derived from X-ray observations.  Moreover, a large fraction ({\\it probably most\\/}) of the cooling flow gas accreted into the emission-line region is converted into stars.  In the light of this new data, we discuss the evolutionary sequence suggested by McNamara (1997), in which radio triggered star formation bursts take place several times during the lifetime of the cooling flow. We conclude that this scenario is consistent with the available observations. ",
        "introduction": "The cooling time of the hot intra cluster medium is shorter than the Hubble time in the central regions of most galaxy clusters (e.g. Arnaud 1988, Edge, Stewart \\& Fabian 1992), which suggests that large amounts of cool material should be settling at the cluster centre.  This process, normally known as a {\\it cooling flow\\/}, may take place at rates of \\mbox{$\\dot{M} \\sim 100$~M$_{\\odot}$~yr$^{-1}$} (Fabian 1994), which implies that \\mbox{$M_{\\rm Total} \\sim 10^{12}$~M$_{\\odot}$} of gas would cool and accrete onto the central cluster elliptical galaxies over the lifetime of the clusters. Considering the idealised picture of an homogeneous cooling flow, all the material will be deposited in the centre of the cluster. However, X-ray data show that the gas is highly inhomogeneous and, very likely, most of the gas must be cooling out at large radii, whilst only a fraction reaches the centre of the flow.  In fact, mass deposition profiles are reasonably fitted by $\\dot{M}(<r) \\propto r$ (Thomas, Fabian \\& Nulsen 1987), although some clusters do not follow this trend beyond a certain radius $r \\ga 60$~kpc (David, Jones \\& Forman 1994; Irwin \\& Sarazin 1995).  For comprehensive reviews on cooling flows see Fabian, Nulsen \\& Canizares (1984, 1991), Sarazin (1986) and Fabian (1994).  A question of open debate is the nature of the final repository of the cool gas. The discovery of significant excess absorption in the soft X-ray spectra of some galaxy clusters containing cooling flows (White et al. 1991; Johnstone et al. 1992; Miyaji et al. 1993; Allen et al. 1993, 1995; Fabian et al. 1994; Irwin \\& Sarazin 1995; Allen \\& Fabian 1997a, 1997b) requires \\mbox{$\\sim 10^{11}$--$10^{12}$~M$_{\\odot}$} of absorbing cold gas, probably in small, pressure-confined, magnetised clouds (Daines, Fabian \\& Thomas 1994). This value is comparable to $M_{\\rm Total}$, the mass deposition rate integrated over the age of the clusters. Fabian, Johnstone \\& Daines (1994) have postulated that dust can form in the cold clouds embedded in cooling flows, which would explain why the clouds are almost undetectable outside the X-ray wavelength range. However, there is a great uncertainty concerning the nature of the cold material producing this absorption (Sarazin 1997, Laor 1997).  Voit \\& Donahue (1995), from the limits imposed by 21~cm and CO observations (McNamara, Bregman \\& O'Connell 1990; Antonucci \\& Barvainis 1994; McNamara \\& Jaffe 1994; O'Dea et al. 1994; O'Dea, Gallimore \\& Baum 1995), suggest that it is unlikely that significant amounts of cold gas can remain undetectable.  On the other hand, very recently Jaffe \\& Bremer (1997), using \\mbox{$K$-band} spectroscopy, have found that the inner few kpc of central cluster galaxies (CCGs) in cooling flows exhibit strong emission in the \\mbox{H$_2$(1--0)S(1) line}, which is not seen in a comparable sample of non cooling flow galaxies.  These authors suggest that it is likely that a large mass of additional molecular material covers the cooling flow as a whole, producing the soft X-ray absorption.  Since it seems clear that at least a fraction of the cool gas must reach the inner regions of the CCGs, it is reasonable to think that part of this gas could accumulate as molecular clouds which, presumably, would form stars.  Unusual diffuse blue light and colours (Sarazin \\& O'Connell 1983; Bertola et al. 1986; Johnstone, Fabian \\& Nulsen 1987 ---hereafter JFN87---; Romanishin 1987; Crawford et al. 1989; McNamara \\& O'Connell 1989, 1992 ---hereafter MO89, MO92---, 1993; Thuan \\& Puschell 1989; N{\\o}rgaard-Nielsen, J{\\o}rgensen \\& Hansen 1990; Allen et al.  1992; Crawford \\& Fabian 1993; Crawford et al.  1995; Allen 1995; McNamara 1995; Hansen, J{\\o}rgensen \\& N{\\o}rgaard-Nielsen 1995; Cardiel, Gorgas \\& Arag\\'{o}n-Salamanca 1995 ---hereafter CGA95---; Melnik, Gopal-Krishna \\& Terlevich 1997), strong low-ionization optical line emission (Baum 1992, Donahue \\& Voit 1997, and references therein), and powerful radio sources (see Jaffe 1992, Burns et al.  1997, for reviews) are usually associated with the inner few kpc of central dominant galaxies in clusters with clear X-ray signatures of cooling flows.  Although with a large scatter, the strength of the optical anomalies are correlated with mass deposition rates (JFN87; MO89; MO92; Heckman et al.  1989 ---hereafter HBvBM---; CGA95; McNamara 1997). In most cases star formation has been invoked to reproduce the excess blue light in the continuum, with typical derived star formation rates ranging from a few to some tens of solar masses per year. It is important to stress that these values normally only account for \\mbox{$\\la 10$} per cent of the total X-ray derived mass deposition rates.  McNamara \\& O'Connell (1993) found that the blue colours observed in A~1795 and A~2597 are located in patches that are coincident with the radio lobes of the central galaxies, which suggests a connection between the radio source and the star formation (or, at least, the source of the blue light). Although power-law spectra could also be an alternative explanation for the origin of the excess blue light (e.g. Crawford \\& Fabian 1993), the absence of polarization in \\mbox{$U$-band} observations of the blue lobes in A~1795 (McNamara et al. 1996b) makes this scenario unlikely. In addition, observations of this object with the {\\it Hubble Space Telescope\\/} (McNamara et al. 1996a; Pinkney et al. 1996), and the {\\it Ultraviolet Imaging Telescope\\/} (Smith et al. 1997), strongly support that the origin of the blue lobes could be an accreted population formed from the cluster cooling flow or from episodes of star formation induced by accretion of stripped material from several gaseous dwarf galaxies.  In this work we address the possible connection between star formation and the cooling flow phenomenon through the study of the stellar populations in CCGs. For this purpose we analyse line-strength gradients in Mg$_2$ (Faber et al. 1985) and the \\mbox{$\\lambda$4000-\\AA\\ break} (D$_{4000}$, definition in Bruzual 1983) in a sample of 11 brightest cluster galaxies.  This paper is an extension of a previous work presented in CGA95 after the inclusion of extensive new spectroscopic data.  The galaxy sample, reduction of the data and error analysis are presented in Sections~2 and~3. The nuclear line-strength and gradients are described in Section~4. The discussion of these results is presented in Section~5, and our conclusions are summarised in Section~6. \\mbox{$H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$} has been assumed throughout the paper. ",
        "conclusions": "\\subsection{Central spectral indices} \\label{sectiondiscussionindex}  Different authors have reported correlations of central blue colours and spectral features with mass deposition rates: JFN87 found a correlation between D$_{4000}$ and the mass flow rate ($\\dot{M}_{\\rm V}$) within their spectrograph slit; MO89, MO92 and McNamara (1997) shown that central $U-B$ colour excesses in CCGs immersed in clusters with high mass deposition rates also correlate with $\\dot{M}$; in CGA95 we confirmed the correlation with D$_{4000}$ and presented a correlation with Mg$_2$.  In Fig.~7 we have plotted all the central measurements listed in Table~3, as a function of the X-ray derived mass deposition rate. In panels~7a and~7b we have used the total $\\dot{M}$ values, whereas in panels~7c and~7d we computed the central mass deposition rates in the same fixed metric aperture size employed in the measurement of the central indices, assuming $\\dot{M} \\propto r$.  We have also included in panels~7a and~7c central D$_{4000}$ estimates of eight additional galaxies obtained from the nuclear \\mbox{(U$-$b)} colours of MO92. For this purpose we have used the linear fit derived in Section~\\ref{secpreviouswork}. These new indices are listed in Table~5.  Filled and open symbols in Fig.~7 correspond to CCGs with and without central emission lines, respectively.  Although, as mentioned above, central D$_{4000}$ and Mg$_2$ indices are well correlated with mass deposition rate in the sense that both indices decrease when the mass flow increases, it is interesting to note that some objects lie off the main relation. In particular the cooling flow galaxies without emission lines A~644, A~2029 and A~2142 are found above the relation, being their central indices fully consistent with those of central galaxies in clusters without cooling flow. Simultaneously, there are also two cooling flow galaxies with emission lines, Hydra~A and N~1275, which also seem to depart from the relation with $\\dot{M}$.  To perform a simultaneous study of the variations in D$_{4000}$ and Mg$_2$, we present an index--index diagram in Fig.~8.  In panel~8a we plot the central indices with their errors for the whole sample.  In the following panels of this figure we have suppressed the galaxy names and error bars for clarity.  Since the CCGs exhibit a clear sequence in this index--index plane, in panel~8b we compare this sequence with that displayed by early-type galaxies. The small dots correspond to a sample of 135 elliptical and S0 galaxies observed by Kimble, Davidsen and Sandage (1989; note: we have transformed their Mgb data into Mg$_2$ using the same relation employed in section~\\ref{secpreviouswork}). As we already shown in CGA95, CCGs in cooling flow clusters populate the D$_{4000}$--Mg$_2$ plane following a different sequence to that displayed by elliptical galaxies.  The enlargement of the galaxy sample allow us to gain new insights into this study. In particular, it is clear from the figure that the CCGs in cluster without cooling flows (open circles) lie in the same region than the giant ellipticals (high Mg$_2$ and D$_{4000}$ values).  The central galaxies in clusters with cooling flows which do not exhibit emission lines in their central parts (open stars, with symbol size proportional to mass deposition rate) are also compatible with the location of the non cooling flow galaxies, even for high $\\dot{M}$, like A~2142 and A~2029. On the other hand, cooling flow galaxies with central emission lines (filled circles, with symbol size proportional to $\\dot{M}$) depart from this location, being the largest accretors the galaxies which show the largest deviations.  Using the same symbols, we plot in panel~8c the D$_{4000}$ obtained from \\mbox{(U$-$b)$_{\\rm nuc}$} colours of MO92 (as explained above). It is interesting to note that some cooling flow galaxies with only moderately high $\\dot{M}$ (Hydra~A and N~1275) display quite low D$_{4000}$ values.  Thus, it is clear that cooling flow galaxies with emission lines do not represent a simple one-parameter sequence in mass deposition rate.  We investigate now the behaviour of the spectral indices as a function of age and metallicity by studying the evolution of single bursts of star formation. We have used the galaxy isochrone synthesis spectral evolution library, GISSEL96, from Bruzual \\& Charlot (1997). In panel~8d we plot the evolution of three single bursts of star formation (full lines) with different metallicities ([Fe/H]=$-$0.64, $-$0.33 and 0.09).  The model sequences are labelled with small circles which indicate the age of the burst in Gyr. Mg$_2$ values in the models are computed using the empirical fitting functions derived by Gorgas et al. (1993) and Worthey et al. (1994).  Unfortunately, this type of fitting functions are not available for the D$_{4000}$, and a non-negligible systematic offset between the predictions of the models and the data could be present.  For simplicity, we represent the sequence of normal ellipticals and S0 of Kimble et al.  (1989) by a dashed parallelogram. Taking into account the models, we have drawn two vectors to indicate changes of age and metallicity in this D$_{4000}$--Mg$_2$ plane and, interestingly, both parameters are not completely degenerated in this diagram. It suggests that normal ellipticals represent a metallicity sequence, whereas the cooling flow objects describe a locus quite parallel to the age evolution. This strongly suggests that the observed spectral properties of the cooling flow galaxies can be understood as a consequence of star formation.  Following the procedure described in Gorgas et al. (1990) and in CGA95, we have modelled the expected changes in the central D$_{4000}$ and Mg$_2$ indices when a normal CCG, initially without cooling flow, undergoes star formation as a result of the cooling flow ---panel~8e. All the models describe the effect of such star formation overimposed on a galaxy with D$_{4000}$ and Mg$_2$ indices equal to the averaged values obtained with our subsample of five CCGs in clusters without cooling flows (i.e., typical of old stellar populations). Dotted lines represent the effect of a continuous star formation during the last 0.1~Gyr, 1.0~Gyr and 10.0~Gyr (from bottom to top). Full lines correspond to single bursts occurred 0.01~Gyr, 0.1~Gyr and 1.0 Gyr ago (also from bottom to top). In all these models a Scalo (1986) IMF and solar metallicity were adopted.  Finally, the dashed line indicates the effect of a power-law spectrum ($F_{\\lambda} \\propto \\lambda^{\\alpha-2}$, with $\\alpha=0.5$). No intrinsic reddening has been taken into account in the models.  Small dots in the model lines indicate the fixed fraction $f_V$ (given by the numbers) of the $V$~light that comes from the accretion population (i.e., the newly-formed stars). In the case of a continuous star formation, this fraction $f_V$ can be related to the SFR using the expression \\begin{equation} \\label{eqcontinuous} {\\it SFR} = \\frac{f_V \\, L_V \\, (M/L)_{\\rm AP}}{t}, \\end{equation} where $L_V$ is the total V~luminosity of the galaxy, $(M/L)_{\\rm AP}$ is the stellar mass-to-light ratio of the accretion population, given by the models, and $t$ is the time over which the star formation has taken place. Similarly, the parametrization for the models involving a single burst of star formation is given by \\begin{equation} \\label{eqburst} {\\it M}_{\\rm AP} = f_V \\, L_V \\, (M/L)_{\\rm AP}, \\end{equation} where ${\\it M}_{\\rm AP}$ is the total mass in stars formed in the accretion population.  Panel~8e reveals that the sequence of cooling flow galaxies does not seem to be very well reproduced with a continuous star formation during the last 1--10~Gyr. Either a {\\it young\\/} ($< 1.0$~Gyr) continuous star formation, or simply a recent burst of star formation $\\sim 0.1$~Gyr ago, can match the observed central indices. In the lower right corner of panel~8e we have also plotted the effect of intrinsic reddening in the measurements. As an example, we represent the corrections for an hypothetical object with D$_{4000}=1.7$ and Mg$_2 = 0.300$ as a function of $E(B-V)$ (the colour excess ranges from 0.0 to 1.0, with ticks spaced 0.1 units). As it could be expected, the effect on Mg$_2$ is almost negligible, whereas the contrary is true for the D$_{4000}$. Although Hansen et al. (1995) have reported \\mbox{$E(B-V) \\leq 0.1$} for Hydra~A, other authors have found non-negligible reddening values ($E(B-V) \\sim$~0.1--0.6) for several cooling flow galaxies (e.g. Hu 1992; Allen et al.  1995;  Allen 1995; Donahue \\& Voit 1993; Voit \\& Donahue 1997). If we allow for the existence of internal reddening, the case for recent bursts instead of a continuous star formation is even stronger.  Nevertheless, it is important to keep in mind that the effect of reddening in the stellar continuum is usually lower than in the emission lines, e.g.  Calzetti (1997) reports $E(B-V)_{\\rm stellar} = 0.44 \\, E(B-V)_{\\rm gas}$ in typical star forming regions.  However, we need to be cautious since, apart from intrinsic reddening, there are additional sources of uncertainty present in the models. First, there is some uncertainty in the starting point of the model sequences, since the dispersion in the central indices of CCGs without cooling flows is probably real.  Second, we have already mentioned that the modelling of D$_{4000}$ is somewhat uncertain given the lack of empirical calibrations, thus some systematic effects could be present. And finally, the models assume solar metallicity {\\it and\\/} solar element ratios: if there is an enhancement of [Mg/Fe] in the cooling flow material, the predicted Mg$_2$ indices could be underestimated. For these reasons, it seems premature to rule out competely, from the models shown in panel~8e alone, a continuous star formation (lasting some Gyrs) as the source of the blue light.  Considering the relatively short wavelength baseline between the D$_{4000}$ and Mg$_2$ indices, the contribution of a power-law spectrum could also reproduce the observed spectral changes. This result is not surprising since, in a short wavelength interval, the spectrum of an early-type star can be approximated quite well by a power law (see also Crawford \\& Fabian 1993). However, as we have already discussed in the introduction, the absence of polarization in the blue light of A~1795 (McNamara et al. 1996b), does not give strong support to this alternative. In addition, some central galaxies show clear signatures of strong Balmer absorption lines (e.g. N~1275, Crawford \\& Fabian 1993; Hydra~A, Hansen et al. 1995, Melnick et al. 1997).  Allen (1995) found that the strong UV/blue continua in the spectra of a small sample of CCGs were better described by young stars than by power-law emission models. This author also reported positive detection of Wolf-Rayet features in two CCGs (those in A~1068 and A~1835), which favour the idea of the formation of massive stars. In this sense, Voit \\& Donahue (1997) have measured a weak He\\,{\\sc ii}~$\\lambda$4686 emission in A~2597. We also find evidences of such helium feature in the central spectrum of the dominant galaxies in A~1795 and in A~2597 (see Fig.~9)  The role of metallicity and the choice of IMF on the model predictions plotted in panel~8e is explored in panel~8f. We show the GISSEL96 (Bruzual \\& Charlot 1997) predictions for different metallicities (as explained in the panel legend), and two IMFs, namely Salpeter (1955) ---thin lines--- and Scalo (1986) ---thick lines---, for the single burst model with an age of 0.1~Gyr. The predictions of the models do not depend strongly on the choice of both [Fe/H] or the IMF.   In order to obtain further constraints on the range of acceptable models, we need to explore a larger wavelength range. We have done so using UV and near-IR data from the literature (see Fig.~10 and description in the figure caption). Although a continuous star formation lasting 10~Gyr could explain the observed optical and near-IR observations, the UV~fluxes are underpredicted.  Recent bursts with ages $\\sim 0.1$--$0.2$~Gyr are simultaneously compatible with both the UV and IR fluxes, and with the measured central line-strength indices. Note however, that at UV wavelengths the effects of dust extinction can be quite dramatic even for moderate colour excesses:  using the interstellar reddening curve of Savage \\& Mathis (1979), and assuming $E(B-V)=0.2$ and $E(B-V)=0.4$ we estimate that the measured UV fluxes need to be corrected upwards by factors of $\\sim 5$--6 and $\\sim 30$--20, respectively. In these cases, is quite obvious that the UV continuum would constrain the upper limit of the burst age to be $< 0.1$~Gyr. This result agrees with the behaviour of D$_{4000}$ and Mg$_2$ with $E(B-V)$ in Fig~8e, indicating that a recent burst would be strongly favoured against constant star formation.   Another piece of evidence which can help to understand the behaviour of the spectral indices of the CCGs comes from the emission line properties. Using line-ratio diagrams HBvBM classified the emission nebulae at the centres of central dominant galaxies into two types. In particular, Type~I corresponds to objects with relatively high [N\\,{\\sc ii}]$\\lambda$6584/H$\\alpha$ ($\\simeq 2.0$) and [S\\,{\\sc ii}]$\\lambda$6717/H$\\alpha$ ($\\simeq 0.7$), whereas Type~II objects exhibit lower line-ratios ([N\\,{\\sc ii}]$\\lambda$6584/H$\\alpha$ $\\simeq 0.9$ and [S\\,{\\sc ii}]$\\lambda$6717/H$\\alpha$ $\\simeq 0.4$). In panels~8g and~8h we have plotted the available HBvBM types for the galaxies in our sample and in MO92 respectively. There is a clear dichotomy in the way Type I (plotted with ``1'' symbols) and Type II (plotted with ``2'' symbols) galaxies populate the central cluster galaxy sequence. This result is not surprising since Type I objects are associated to clusters with low X-ray and H$\\alpha$ luminosities, small optical nebulae and low mass deposition rate, whereas the contrary is true for Type II objects (HBvBM, Baum 1992, Donahue \\& Voit 1997).  Another clue comes from the morphology of the blue excess light. Recently, McNamara (1997) has established a morphological classification based on the study of the surface brightness distributions and colour structure in cooling flow CCGs. This classification consists in four morphological types ordered by increasing geometrical complexity: (1) point source (e.g. A~2199, A~2052) with unresolved blue nuclei and nebular line emission, low nebular luminosity, with modest blue anomalies;  (2) disk (e.g.  Hydra~A) unusual in cooling flow galaxies and characterised by a disk of gas and young stars rotating around the nucleus; (3) lobe (e.g.  A~1795, A~2597) with bright, blue lobes of optical continuum separated several kiloparsecs from the nucleus; and (4) irregular-amorphous (e.g. N~1275 may be in transition between types~3 and~4). In Fig.~8g we have also plotted the four morphological types in the neighbourhood of the galaxies which display such characteristics.  Taking into account the D$_{4000}$ exhibited by N~1275, we have tentatively located Type~4 near PKS$0745-191$. Very interestingly, McNamara (1997) has speculated about the possibility that the structural types reflect actually an evolutionary sequence. We will discuss this point in more detail in section~\\ref{section:wholepicture}.  \\subsection{Line-strength gradients} \\label{sectiondiscussiongradients}  Although line-strength gradients in Mg$_2$ (e.g. Faber 1977; Gorgas et al. 1990; Gonz\\'{a}lez 1993; Davies et al. 1993; Carollo et al. 1993; FFI95; Gonz\\'{a}lez \\& Gorgas 1995) and in D$_{4000}$ (e.g. Munn 1992; Davidge \\& Grinder 1995) are usually found to be linear with $\\log(r)$ in elliptical galaxies, in CGA95 we showed that some cooling flow galaxies exhibit a clear slope change at intermediate radii. With the inclusion of the new data presented in this paper (see Figs.~3 and~4), we find that this change of slope is clearly related to the presence of emission lines in the central regions of CCGs.  We measured the radius at which the change of slope takes place using error-weighted least-squares fits to two straight lines forced to join at a variable radius. The method gives the {\\it break radius\\/} which minimises the residual variance of the fits.  We estimated the uncertainty in the break radius using Monte Carlo simulations that take into account the measurement errors in the indices (typically $\\simeq 1000$ simulations per gradient). For the emission line galaxies with reasonably good data beyond the emission region, we have compared in Fig.~11 the break radius with the observed emission-line radius (as defined in section~4.2).   The figure reveals that the break radii are almost coincident with the emission-line radii, with the former being slightly larger.  It is quite evident from Figs.~3 and~4 (see also the fitted gradients tabulated in Table~4), that cooling flow galaxies {\\it with emission lines} exhibit flat and even positive gradients in the region where the emission is detected (e.g. A~478, A~1795, A~2597, Hydra~A, A~85, A~2199). In fact, these inner gradients (measured in the emission line region) seem to be correlated with the mass deposition rate, i.e. more positive with increasing $\\dot{M}$ (see Figs.~12a and~12b). In the outer parts of these galaxies, where emission is not observed, the derived mean D$_{4000}$ gradient is consistent with that observed in normal elliptical galaxies (see Fig.~12c). However, the mean Mg$_2$ gradient in the outer regions of cooling flow galaxies with emission lines seems to be marginally stepper than the mean gradient for normal ellipticals (Fig.~12d), although the quality of the Mg$_2$ data at large galactocentric distances is not good enough to reach a firm conclusion.  A couple of galaxies (A~1126 and A~262) seem to have slightly stepper gradients in the inner parts than in the outer regions, although this result is not very significant due to the poor S/N in the outer spectra.  In addition, mean gradients in central dominant galaxies with cooling flow but {\\it without emission lines\\/} (e.g.  A~644, A~2142) are similar to those in galaxies without cooling flows and in elliptical galaxies.  If we assume that the observed slope changes in the line-strength gradients are due to radial index variations induced by star formation in the emission-line region, we can use that information, coupled with the star formation models, to estimate the amount and spatial profile of the accretion population. We assume that the gradients determined in the outer regions represent the {\\it underlying\\/} gradient, i.e., that of the galaxy before the cooling-flow induced star formation takes place. That is justified by the similarity of these outer gradients with those of ``quiescent'' galaxies (i.e., both normal giant ellipticals and CCGs without cooling flows).  In this analysis we have exclusively concentrated on the D$_{4000}$ gradients, which  are much better determined than the Mg$_2$ ones.  As discussed in section~5.1, a single burst model with an age $\\simeq 0.1\\,$Gyr, solar metallicity and Scalo IMF is able to reproduce the observed spectral properties of the cooling flow galaxy sequence. We have employed this particular model to tentatively obtain the spatial profile of the accretion population.  First we computed the averaged indices inside the emission line region (simulating a circular aperture by weighting with light and radius), both from the fitted inner gradient and from the extrapolation of the outer gradient into the emission line region.  Then, we used the star formation model to derive the mean $f_V$ of the accretion population required to reproduce the observed changes in the indices (i.e., the difference between the measured values and those estimated from the extrapolation of the outer gradient). Assuming that the absolute magnitude corresponding to the old stellar population in the emission line region is $\\langle M_V \\rangle \\sim -22.3$~mag\\footnote{Hoessel, Gunn \\& Thuan (1980) reported that the mean magnitude of a sample of brightest cluster galaxies inside a metric radius of $16/h_{60}$~kpc is $\\langle M_V \\rangle = -22.68 \\pm 0.03$~mag. In addition, Hoessel (1980) showed that the internal regions of brightest cluster galaxies are adequately described by a modified Hubble law of the form $I(r)=I_c/(1+r^2/r_c^2)$, where $I_c$ is the central intensity, $r_c$ the core radius and $r$ the angular distance from the centre. Assuming this intensity profile, the integrated luminosity inside a radius $r$ is then given by $L(r)=\\pi \\, I_c \\, r_c^2 \\, \\ln(1+r^2/r_c^2)$. Introducing the value of the Hubble constant adopted in this paper (\\mbox{$H=50 \\; {\\rm km} \\;  {\\rm s}^{-1} \\; {\\rm Mpc}^{-1}$}), and the mean core radius derived by Hoessel (1980), $\\langle r_c \\rangle = 2.23/h_{60}$~kpc, we obtain \\mbox{$\\langle M_V \\rangle \\sim -22.3$~mag} inside a radius \\mbox{$r\\sim 11/h_{50}$~kpc}, which is the average radius of the emission line regions in our sample.}, the mean $f_V$ value is then employed to compute the total $V$~luminosity ($L_V$) of the galaxy plus the burst. Using Eq.~\\ref{eqburst}, and the mass-to-light ratio of the accretion population given by the model, we obtain the total mass in stars of the accretion population formed in the last burst. Note that, since the excess blue light is completely dominated by the youngest stars, this method is only sensitive to the stars formed in the most recent burst.   In Fig.~13 we have plotted the estimated total amount of mass transformed into stars in the emission-line region versus the mass deposition rate computed in the same area (assuming that \\mbox{$\\dot{M} \\propto r$}). There is an increase in the mass transformed into stars with $\\dot{M}$, with Hydra~A lying somewhat outside this trend.  Interestingly, the total mass accreted by the cooling flow in a $\\sim0.1\\,$Gyr period inside the emission line region is remarkably similar to the mass in new stars (see section~\\ref{section:wholepicture}).  Assuming spherical geometry, it is possible to apply the same technique to obtain the amount of mass transformed into stars in concentric shells at different radii. Using a deprojection algorithm described in Appendix~A, we have derived the spatial profiles for the density of mass transformed into stars (Fig.~14).  We also show (dashed line) the expected density profile of the mass deposition assuming $\\dot{M}(<r) \\propto r$.  Interestingly, the derived profiles seem to be quite parallel to the dashed line (the mean slope for all the galaxies excluding Hydra~A is $-$2.08,~\\mbox{${\\rm rms} = 0.18$}), whereas Hydra~A exhibits a statistically significant steeper trend.  This result strongly suggests a direct connection between the mass deposition and the star formation profiles. Indeed, figures~13 and~14 suggest a very close link between the star formation and the cooling flow phenomenon. Such close connection would be difficult to understand without most of the gas accreted in the emission-line region being transformed into stars."
    },
    "9804/astro-ph9804130_arXiv.txt": {
        "abstract": "The light curves of the two images of the double quasar \\object{Q0957+561} as obtained by Kundi\\'c et al. (1997) are almost identical, except for an overall time delay and scaling factor.  This allows us to put limits on the amount of microlensing that took place during the time interval corresponding to the monitoring observations. We perform numerical simulations in which we model the microlensing behaviour of the (halo of the) lensing galaxy in the system.  We test ``MACHO-masses'' ranging from $10^{-8}$ to $10^{-1}$ M$_{\\odot}$ and quasar sizes from $10^{14}$ to $3\\times 10^{15}$ cm. Statistically comparing the expected microlensing-induced changes from 100\\,000 simulated light curves over a period of 160 days with the (lack of) observed fluctuations, we can constrain regions in the parameter space of MACHO mass and quasar size with various degrees of confidence.  In particular, a halo consisting of objects at the low end of our mass scale can be ruled out with high confidence for a small quasar size.  A halo consisting of objects with $10^{-2}$ or $10^{-1}$ M$_{\\odot}$ cannot be ruled out yet, but it should produce MACHO induced fluctuations in future observations. We also test halos with only 50\\% or 25\\% of the mass in compact objects; constraints here are a bit less stringent. ",
        "introduction": "It has been known for almost three decades that rotation curves of galaxies remain flat even beyond the visible matter (Rubin \\& Ford \\cite{Rubin1970}, Rubin et al. \\cite{Rubin1985}). As an explanation, Ostriker et al. (\\cite{Ostriker1974}) and Einasto et al. (\\cite{Einasto1974}) suggested that galaxies are surrounded by large halos consisting of ``dark matter'' and dominating the total mass of the galaxies. Over the years a very large number of possible ``candidates'' for this dark matter have been suggested (for recent reviews see, e.g., Bahcall \\cite{Bahcall1997} or Raffelt \\cite{Raffelt1997}). They can broadly be divided into ``elementary particle'' candidates (e.g. massive neutrinos or axions) and ``astrophysical'' candidates (e.g. black holes, brown dwarfs, ``Jupiters'', comets). Not much progress has been made to date in identifying the elusive dark matter despite major efforts in many directions. In fact, a new branch of physics established itself -- ``astro-particle physics'' -- whose major goal is the solution of the dark matter problem.  More than 10 years ago Paczy\\'nski (\\cite{Paczynski1986}) proposed a direct test to prove or reject the possibility that the halo of the Milky Way consists of dark compact objects by way of gravitational microlensing: If the brightness of at least a million stars in the Large Magellanic Cloud (LMC) could be regularly observed, at any given time about one of them should be magnified due to a halo object (in the following referred to as ``MACHO'', for ``massive compact halo object'') passing in front of it and focussing the light rays to the observer. The light curve of the affected background star should show a very characteristic and achromatic behaviour.  Soon thereafter various groups started big observational programs to investigate this promising possibility, and the first microlensing events were found in 1993 (Alcock et al. \\cite{Alcock1993}; Aubourg et al. \\cite{Aubourg1993}). The most recent results indicate that about $50^{+30}_{-20}\\%$ of the dark matter in the Milky Way halo can consist of such objects (Alcock et al.  \\cite{Alcock1997}). Their most likely mass range ($M \\approx 0.5 ^{+0.3}_{-0.2}$ M$_\\odot$), however, is higher than originally suggested for genuine ``brown dwarfs'', but the uncertainty is large.  The largest problem in these experiments is the small probability for a microlensing event: the ``optical depth'' (the fraction of background stars that are significantly affected by microlensing) is less than $10^{-6}$.  In a very different optical depth regime, gravitational lensing can also be used to test whether halos of {\\it other} galaxies are made of compact objects (Gott \\cite{Gott1981}): In cases of good projected alignment along the line of sight between a background qua\\-sar and a foreground galaxy, the galaxy can produce multiple images of the quasar\\footnote{Note that here the distances involved are about five orders of magnitude larger than in the microlensing searches towards the LMC (few Gigaparsec rather than some 55 kpc), and the optical depths involved are of order unity (i.e. six orders of magnitude higher).}. In this case the light bundles from the quasar pass through (the dark halo of) the foreground galaxy on their way to the observer. These light bundles probe the graininess of the halo: if the dark matter there consists of some kind of elementary particles (``smoothly distributed matter'' on astronomical length scales), the light bundles should be unaffected; if it is made of compact astrophysical objects, the light bundles can probe it.  The measured brightnesses of the quasar images should vary as a function of time due to the changing relative positions between lens, source and observer (Chang \\& Refsdal \\cite{Chang1979}).  The surface mass density (optical depth) in these cases is high enough (of order unity) to basically affect (change) the measured flux of a quasar image all the time.  Since quasars are intrinsically variable as well, it is not trivial to decide whether an observed variability is intrinsic or microlensing-induced. In the case of multiple images of one quasar, however, the intrinsic fluctuations should show up coherently in all the images -- modulo a certain time delay -- whereas the microlensing changes occur uncorrelated in the various images. Once well-sampled light curves of two (or more) quasar images are obtained and the time delay $\\Delta t$ (due to the different light paths of different light bundles) and the magnitude difference $\\Delta m$ (due to different magnifications) is known (both of which are not trivial), one can shift the two light curves in time and magnitude by the appropriate amounts $\\Delta t$ and $\\Delta m$ and subtract them from each other.  All remaining fluctuations in the ``difference light curve'' must be due to microlensing. In particular, if the difference light curve is flat, this indicates that there was no microlensing going on during the period of observation\\footnote{The effect of ``flat'' microlensing -- periods in which microlensing produces a constant (de-) magnification -- are taken care of by a different magnification ratio.}.  The gravitationally lensed double quasar Q0957\\linebreak[0]+561 (Walsh et al. \\cite{Walsh1979}) has been observed/monitored for almost two decades by many groups (see e.g., Schild \\& Thomson \\cite{Schild1995}) with the goal to measure the time delay and subsequently determine the Hubble constant.  The time delay is now established firmly at $\\Delta t = 417\\pm3$ days (Schild \\& Thomson \\cite{Schild1997}; Kundi\\'c et al. \\cite{Kundic1997}, subsequently K97; Oscoz et al. \\cite{Oscoz1997}). In addition, a number of gravitational lens models have been published for the Q0957+561 system which reproduce the observed (optical and radio) image configurations well (Falco et al. \\cite{Falco1991}; Grogin \\& Narayan \\cite{Grogin1996}), so that the parameters relevant for microlensing at the position of the quasar images can be determined from these models with an accuracy of a few percent.  With both the time delay established and good models being available, the double quasar Q0957+561 can hence be used as a test for massive halo objects by comparison of the light curves of images A and B.  We do this here by simulating the microlensing effect numerically for MACHOs of different masses.  We analyse the resulting microlensed light curves and show how often microlensing-induced changes of certain amplitudes are to be expected for certain MACHO masses. Finally, we compare our numerical simulations with the well-sampled data set of K97 and thus constrain the MACHO masses in (the halo of) the lensing galaxy.  In Sect. 2 we describe briefly the monitoring data set of K97 for the double quasar and the simulations we perform in order to compare the observations with microlensing light curves due to different masses.  In Sect. 3 we show our results for a MACHO mass range from $10^{-7} < M/M_{\\odot} < 10^{-1}$, and in the final Sect. 4 we present our conclusions. ",
        "conclusions": "In two years of observational data on the gravitationally lensed double quasar Q0957+561 by Kundi\\'c et al. (\\cite{Kundic1997}) no microlensing variation of the quasar larger than $0.05$ magnitudes was observed. In our microlensing simulations we find that this rules out a dominant population of compact objects in the halo of the lensing galaxy from $10^{-3}\\,M_{\\odot}$ down to $10^{-5}\\,M_{\\odot}$ for quasar sizes below $3\\times 10^{14}\\,{\\rm cm}$. These limits are almost independent of the fractions of the halo mass contained in compact objects.  These results are consistent with the 2nd year results from the MACHO microlensing search towards the Magellanic clouds (Alcock et al. \\cite{Alcock1997}). The most probable mass of MACHOs in the Milky Way that emerges from their study lies around half a solar mass, which is not excluded by our study. In fact, such large masses are not yet probed by the observational data set we used (K97) because microlensing effects by objects with masses of about one solar mass are only becoming observable on time-scales of several years. In a different data set on Q0957+561 (Schild \\cite{Schild1996a}) there were indications for some microlensing action which can be produced by stars in this mass range. Very recently, Pelt et al. (\\cite{Pelt1998}) analysed a large data set on Q0957+561. Similar to us, they cannot find evidence for microlensing variations on short time scales.  Along the same lines, our results are consistent with the results of Lewis \\& Irwin (\\cite{Lewis1996}) who found MACHO masses to be in the range from $0.1 M_{\\odot}$ up to $10.0$ $M_{\\odot}$ from microlensing in the quasar Q2237+0305.  The limits on compact objects in the halo of the lensing galaxy of Q0957+561 will, however, improve with time since the monitoring of Q0957+561 at the Apache Point Observatory (and also at other observatories) is an ongoing project. If, for example, microlensing variations above $0.05$ magnitudes are not observed for another season, the limits from Tables~\\ref{macho1}, \\ref{macho2} and~\\ref{macho3} would become stronger and we could possibly rule out compact objects with masses as high as $0.1 M_{\\odot}$ or more.  In quasar microlensing, not only the mass of MACHOs, but also the size of the quasar enters the calculations. With the current data, however, we cannot constrain the quasar size very much. To do this, one would need to find characteristic events in the difference light curve of the two quasar images. If, for example, the peak in the difference light curve (Schild \\cite{Schild1996b}, K97) in Fig.~\\ref{diff} around day 750 is real, it would certainly be a valuable constraint for both the masses of MACHOs and the quasar size.  Finally, in magnification patterns with side lengths of several hundred Einstein radii we find coherent structures in the caustic network on scales of 200 Einstein radii or more. In the lensing galaxy of Q0957+561 these scales correspond to physical scales of the order of $1.7 \\sqrt {\\frac{M}{M_{\\odot}}}\\,h^{-1/2}$pc for MACHOs of mass $M$. On these scales the dark objects might not be distributed completely randomly anymore. It would thus be interesting to investigate further the effect of clustering of MACHOs on the microlensing properties of galaxies."
    },
    "9804/astro-ph9804306_arXiv.txt": {
        "abstract": "X-ray temperature measurements of clusters of galaxies are now reaching to redshifts high enough to constrain $\\Omega_0$. A redshift-dependent relation that maps these X-ray temperatures to the virial masses of clusters is an essential ingredient when one is trying to determine cosmological parameters from cluster evolution. Most such relations assume that clusters form from top-hat perturbations that virialized just before the time we are observing them.  The smaller $\\Omega_0$ is, the less accurate these relations become.  Here we derive a relation between virial mass and cluster temperature that allows for the fact that clusters form gradually and cease forming when the density of the universe drops well below the critical value.  We show how sensitively the expected redshift distribution of clusters depends on the mass-temperature relation used and argue that one needs to use a relation that yields no evolution in the low-$\\Omega_0$ limit. ",
        "introduction": "Studies of temperature evolution in massive clusters of galaxies are now realizing their promise as cosmological indicators. The most massive clusters in the universe mark the rarest, highest amplitude peaks in the intial perturbation spectrum that seeded structure formation and provide powerful leverage on parameters such as $\\Omega_0$, the current matter density of the universe in units of the critical density $\\rhocr$.  If $\\Omega_0 = 1$, the development of structure should have proceeded unabated from early in time to the present day. However, the relative lack of evolution observed in the X-ray temperature function of clusters indicates that the formation of structure is stalling, and thus that $\\Omega_0 \\approx 0.2-0.5$ (Henry 1997; Bahcall \\etal 1997; Carlberg \\etal 1997; Donahue \\etal 1998)  Using X-ray observations of clusters to constrain cosmological models requires knowledge of how the X-ray temperature of a cluster ($\\tx$) relates to its mass.  Temperatures of distant clusters are now straightforward to measure (e.g., Donahue 1997; Donahue \\etal 1998; Henry 1997), but relating these temperatures to cluster masses is a subtler art. In order to define a cluster's mass, one must decide upon an appropriate boundary separating the cluster itself from peripheral material not yet part of the cluster. Theorists tend to choose a boundary that encompasses matter that has virialized and excludes matter that is still falling into the cluster.  The mass within this boundary is then $\\mvir$, the virial mass of the cluster.  Defining a cluster's mass to be $\\mvir$ has an established theoretical pedigree, but this definition is observationally problematic. Virial masses are a staple of the Press-Schechter approach to the problem of hierarchical structure formation, which has become enormously popular owing to its basic agreement with much more sophisticated $n$-body techniques (e.g., Lacey \\& Cole 1994). Press \\& Schechter (1974) presumed that the collapse of a spherical top-hat perturbation would adequately represent the formation of more complicated objects.  In such a spherically symmetric collapse, the density of a newly formed object just after collapse is $18 \\pi^2 \\rhocr \\approx 178 \\rhocr$ in an $\\Omega_0 = 1$ universe, but this contrast factor drops to as low as $8 \\pi^2 \\rhocr \\approx 79 \\rhocr$ in an open universe with $\\Omega_0 \\ll 1$ (e.g., Lacey \\& Cole 1993).  Even if one could measure such density contrasts accurately, {\\em a priori} knowledge of $\\Omega_0$ would still be needed to determine $\\mvir$. Thus, the bridge between $\\tx$ and $\\mvir$ must be built primarily from the theoretical side.  Numerical models of cluster formation show that cluster temperatures should indeed be closely related to their masses (e.g., Evrard \\etal 1996). To lowest order, cluster potentials that arise in such models are isothermal, with $T_X \\propto M(r) / r$ (e.g., Cole \\& Lacey 1996). If the mean density within radius $r$ is $\\Delta$ times the critical density, then the mass within that radius is $M(\\Delta) \\propto T_X^{3/2} \\rhocr^{-1/2} \\Delta^{-1/2}$. Proceeding in the vein of the Press-Schechter approach, one can define $\\Delta_{\\rm vir}$ to be the density contrast of a spherical top-hat perturbation just after collapse and virialization, in which case $\\mvir = M(\\Delta_{\\rm vir})$ and $T_X \\propto \\mvir^{2/3} \\rhocr^{1/3} \\Delta_{\\rm vir}^{1/3}$.  The density contrast parameter $\\Delta_{\\rm vir}$ depends in general on $\\Omega_0$, the redshift of virialization, and the value of the cosmological constant (e.g., Eke \\etal 1996; Oukbir \\& Blanchard 1997). Adopting the recent-formation approximation, which assumes that the clusters we now observe virialized just before their light departed, then leads to a relation like the following: \\begin{equation} \\label{ecf} k T_X = (1.38 \\, {\\rm keV}) \\beta^{-1} \\left( \\frac {\\mvir} {10^{15} h^{-1} M_\\odot} \\right)^{2/3} \\left[ \\frac {\\Omega_0} {\\Omega(z)} \\right]^{1/3} \\Delta_{\\rm vir}^{1/3} (1+z) \\; \\; , \\end{equation} where $\\Omega(z)$ is the matter density of the universe in units of $\\rhocr$ at redshift $z$, $\\beta$ is the usual ratio of potential depth to X-ray temperature, and $h$ is the Hubble constant in units of $100 \\, {\\rm km \\, s^{-1} \\, Mpc^{-1}}$ (e.g., Eke \\etal 1996). In the limit of $\\Omega(z) \\approx 1$, we arrive at the scaling relation $T_X \\propto \\mvir^{2/3} (1+z)$.  The shortcomings of this approach are well known (e.g. Viana \\& Liddle 1996; Kitayama \\& Suto 1996; Eke \\etal 1996).  Real clusters do not form at a single moment but rather accrete their matter over a long period of time.  If the accretion rate remains sufficiently high, as would be the case if $\\Omega_0 = 1$, then the clusters we now see attained their observed masses relatively recently, validating the recent-formation approximation that led to equation (\\ref{ecf}).  If instead $\\Omega_0 \\ll 1$, then the recent formation approximation grows increasingly less valid in time.  Note that in the limit $\\Omega_0 \\rightarrow 0$, cluster evolution ceases by $z = 0$ and the $\\mvirtx$ relation should no longer depend on redshift, whereas equation (\\ref{ecf}) dictates $T_X \\propto M_{\\rm vir}^{2/3} (1+z)^{2/3}$.  One should also keep in mind that the temperature of a real cluster depends on formation epoch only insofar as the formation epoch affects the shape of the halo density profile and furthermore that real clusters are not necessarily isothermal. Comparisons of simulated clusters with equation (\\ref{ecf}) indicate that the effective value of $\\beta$ declines modestly but systematically with time in an open universe with $\\Omega_0 = 0.3$, slightly diminishing the expected amount of temperature evolution (Eke, Navarro, \\& Frenk 1998).  While the relation in equation (\\ref{ecf}) is adequate for many purposes, one would prefer to use a relation valid in the low-$\\Omega_0$ limit when constraining $\\Omega_0$ through cluster temperature evolution. One way to patch the recent-formation approximation is to define an explicit formation redshift $z_f$ at which a cluster virializes. One can then integrate over the appropriate distribution of formation redshifts to determine the properties of observed clusters at redshift $z$ (Kitayama \\& Suto 1996; Viana \\& Liddle 1996).  Motivated by the fact that clusters form gradually, not instantaneously, we present here an alternative approach to characterizing the redshift evolution of the $\\mvirtx$ relation which has the advantage of being qualitatively correct in the low-$\\Omega_0$ limit.  In \\S~2 we idealize clusters as forming from spherically symmetric perturbations with negative radial density gradients rather than from top-hat perturbations. This approximation alleviates the need to specify an artificial formation redshift because it explicitly accounts for how clusters grow with time. It also yields an analytical formula for the change in virial energy with virial mass.  We then use the merging-halo formalism of Lacey \\& Cole (1993) to derive an expression for the time-averaged mass accretion rate of a massive cluster, which allows us to determine how the $\\mvirtx$ relation should change with time.  Section~3 briefly illustrates some implications of this relation. ",
        "conclusions": ""
    },
    "9804/astro-ph9804261_arXiv.txt": {
        "abstract": "The discovery of proto--globular cluster candidates in many current--day mergers allows us to better understand the possible effects of a merger event on the globular cluster system of a galaxy, and to foresee the properties of the end--product. By comparing these expectations to the properties of globular cluster systems of today's elliptical galaxies we can constrain merger models. The observational data indicate that {\\it i}) every gaseous merger induces the formation of new star clusters, {\\it ii}) the number of new clusters formed in such a merger increases with the gas content of the progenitor galaxies. Low--luminosity (about $M_V>-21$), disky ellipticals are generally thought to be the result of a gaseous merger. As such, new globular clusters are expected to form but have not been detected to date. We investigate various reasons for the non--detection of sub--populations in low--luminosity ellipticals, i.e.~absence of an old population, absence of a new population, destruction of one of the populations, and finally, an age--metallicity conspiracy that allows old and new globular clusters to appear indistinguishable at the present epoch. All of these possibilities lead us to a similar conclusion, namely that low--luminosity ellipticals did not form recently ($z<1$) in a gas--rich merger, and might not have formed in a major merger of stellar systems at all. High--luminosity ellipticals do reveal globular cluster sub--populations. However, it is difficult to account for the two populations in terms of mergers alone, and in particular, we can rule out scenarios in which the second sub--population is the product of a recent, gas--poor merger. ",
        "introduction": "Globular clusters are thought to be good tracers of the past evolution of their parent galaxy. Recreating the history of a galaxy from the current properties of its globular cluster system is complex and challenging. However, if a specific event is expected in the history of a galaxy, such as a major merger, and this event is supposed to leave a clear signature in the globular cluster system (e.g.~the formation of a new population of globular clusters), then studies of globular cluster systems can be used to confirm or rule out the occurrence of such events.  In this paper, we examine the constraints that can be set from the current observations of globular cluster systems on the merger model for elliptical galaxies. The main challenge is to understand the apparent presence of {\\it only one} population of globular clusters in the luminosity and colour distributions of small ellipticals. Ashman \\& Zepf (1992) made a variety of predictions for the properties of globular cluster systems after a merger of two gas--rich galaxies. One of their main predictions is that mergers will leave a clear signature in the form of a {\\it second} population of globular clusters, formed during the merger,  which is added to the old populations of globular clusters brought in by the progenitors. A notable success of their model is the presence of young proto--globular cluster candidates in currently merging systems, confirming the basic idea that such events can modify the existing globular cluster systems. Furthermore some of their predictions seem to be verified for the globular cluster systems of giant elliptical galaxies (Zepf \\& Ashman 1993): two or more globular cluster populations are detected in colour distributions, the red (presumably new) globular clusters are more concentrated toward the center of the galaxy as expected if these formed from the in--falling gas.  However, whether the globular cluster systems of all ellipticals verify these predictions remains an open question. The properties of globular clusters in small and large ellipticals seem to differ. Studies of systems in small ellipticals failed to detect more than one population of globular clusters (Kissler-Patig 1997a). Moreover, Forbes, Brodie \\& Grillmair (1997) have recently argued that the observational data on globular clusters in high--luminosity ellipticals (i.e.~gE and cD galaxies) cannot be explained in detail by an Ashman \\& Zepf type merger, and could be explained in other formation scenarios.  We adopt a working definition of --18 $>$ M$_V$ $>$ --21 for low--luminosity and --21 $>$ M$_V$ $>$ --23.5 for high--luminosity ellipticals (for H$_0$ = 75 km s$^{-1}$ Mpc$^{-1}$) throughout this paper. Further, we will refer to the {\\it new} population of globular clusters as the population formed in the merger event, and to the {\\it old} population of globular clusters as the ``original'' population, present before any merger event.  We examine the numerous observations of proto--globular clusters in merging galaxies to extract the signature that these events will leave on the globular cluster systems (Sect.~2). In particular, we examine if, and how many, globular clusters are expected to form in the different (gas--rich/gas--poor) mergers. In Section 3, we consider low--luminosity and high--luminosity galaxies in turn. We investigate the various alternatives that could explain the presence of only one population of globular clusters in low--luminosity ellipticals, despite the fact that they are generally thought to have formed several Gyr after a gaseous merger event and thus should show a second population of globular clusters. The most interesting alternative (Sect.~3.2) is the presence of two populations whose mean ages and metallicities conspire to let them appear the same in colour and magnitude distributions. The various alternatives allow us to put constraints on the occurrence of a merger event in the recent history of the parent galaxy. We then examine in Sect.~3.4 if similar arguments can also constrain the history of high--luminosity ellipticals, and in Sect.~3.5 if additional constraints can be obtained from the properties of ongoing mergers. We summarize our findings and conclusions in Sect.~4. ",
        "conclusions": ""
    },
    "9804/astro-ph9804057_arXiv.txt": {
        "abstract": "We study linear power spectra and formation of large scale structures in flat cosmological models with $\\Lambda \\ge 0$ and cold plus hot dark matter components. We refer to these models as mixed $\\Lambda$ models (MLM). The hot component consists of massive neutrinos with cosmological density $\\Omega_H$ and number of neutrino species as a free parameter. The linearized Einstein-Boltzmann equations for the evolution of the metric and density perturbations are integrated for a set of values of the cosmological parameters. We  study MLM models with present matter density in the range $ 0.25 \\le \\Omega_M \\le 1$, dimensionless Hubble constant $0.4 \\le h\\le 0.7$ and the hot dark matter content with a ratio within the limits $ 0 \\le \\Omega_H/\\Omega_M\\le0.3$. For all the considered models we assume a scale-invariant primeval spectrum.  The density weighted final linear power spectra are normalized to the four year COBE data and have been used to constrain the parameter space  by a comparison of  linear predictions with the current observational data on large scales. The consistency of MLM predictions with the observable data set is best obtained for models with one species of massive neutrinos and $\\Omega_H/\\Omega_M\\le0.2$. Of the considered linear tests the strongest constraints on $\\Omega_M$ that we obtain arise by comparing the cluster X-ray temperature function with that observed at the present epoch. Consistency with the estimated cluster abundance can be achieved for COBE normalized MLM models with $\\Omega_H/\\Omega_M\\le0.2$ and $ 0.45 \\le \\Omega_M \\le 0.75$ for $h=0.5$. If $h=0.7$ then $ 0.3 \\le \\Omega_M \\le 0.5$. These constraints are at $1\\sigma$ level and standard MDM models are clearly ruled out.  We note that the range of allowed values for $\\Omega_M$, that we obtain for MLM models from linear analysis, is also approximately the same range that is needed in order to consistently satisfy a variety of independent observational constraints. ",
        "introduction": "In the standard framework of gravitational instability theory present day structures must have been formed through the growth of small inhomogeneities from an initial random Gaussian density field present at very early epochs, with a scale-invariant Harrison-Zel'dovich spectrum. Clustering analysis of the large scale structure in the Universe has been improved in recent years by observations of the spatial distribution of galaxies and cluster of galaxies, as well as cosmic microwave background (CMB) anisotropies. Thus any theory that wants to fit the observed large scale clustering must be consistent with a set of constraints over more than three decades in length: from galaxy correlation ($\\sim 1 {h}^{-1} Mpc$, $H_0=100h Km sec^{-1} Mpc^{-1}$ ) up to the quadrupole CMB anisotropies detected by COBE ($\\sim 3000 {h}^{-1} Mpc$, \\cite{smo92}).  In the standard FRW  metric the fundamental background cosmological parameters are related by a mutual relation. It is also understood that for these parameters (the present matter density $\\Omega_M$, the value of the Hubble constant $H_0$ and the age of the universe $t_0$)  the range of values allowed by observations must be consistent with their FRW relation. If we adopt the standard inflation theory we shall assume that the total energy density in the Universe is equal to the critical density ($\\Omega_0=1$). On the other hand, observations and theoretical predictions suggest that the amount of baryon density must be small ($\\Omega_b \\sim 0.05 {(h/0.5)}^{-2}$, \\cite{wal91}, \\cite{cop95}). A much larger contribution to the matter density must be of non-baryonic nature. This is so-called dark matter (DM) problem and it has observational support from dynamical estimates ($\\Omega_M \\sim {0.2 - 0.3}$, \\cite{bah96} and references cited therein). It must be stressed that there is not yet a firm evidence for $\\Omega_M=1$, this value for $\\Omega_M$ being required by the inflationary paradigm. The simplest possible model of DM is the one in which the universe is dynamically dominated by a single massive collisionless particle. The most popular DM candidates are collisionless massive particles, which decoupled from cosmological plasma either when being relativistic (DM particles like massive neutrinos -- Hot DM) or non-relativistic (hypothetical massive particles -- Cold DM).  The simplest model is CDM, where the power spectrum of the perturbations depends on a single parameter, the present cold density $\\Omega_C$. Historically the first model to be considered was a neutrino with a non-zero rest mass HDM (\\cite{bon83}; \\cite{zak80}). The HDM model was soon rejected because numerical simulations ( \\cite{whi83}) produced nonlinear structures too late to be in agreement with QSO existence. The standard CDM model (SCDM) has been analysed in considerable details (\\cite{dav85}; \\cite {dav88}  and references cited therein). It can match galaxy clustering ($\\mincir 10 {h}^{-1} Mpc$) with a bias parameter $b_g \\magcir 2$, although it lacks of sufficient power on large scales. The SCDM model however became seriously inconsistent with clustering data after the COBE detection of a quadrupole anisotropy in the CMB (\\cite{smo92}). When the power spectrum of the model is normalized to the COBE data SCDM models have unavoidable difficulties (\\cite{wha93}; \\cite{oli93}; \\cite{moa93}; Jing \\& Valdarnini 1993; Dalton et al. 1991 ; Baugh \\& Efstathiou 1994).  The Mixed Dark Matter models (MDM) was  at first proposed~and~discussed~in a few papers (Fang, Xiang \\& Li 1984; Shafi \\& Stecker 1984; Achilli, Occhionero \\& Scaramella 1985; Valdarnini \\& Bonometto 1985) as an example to overcome the standard Hot Dark Matter (SHDM). Later, SMDM models were readressed again (\\cite{hol89}; \\cite{van92}). Finally, it was proposed that a certain mixture proportion of about $30/70\\%$ for the Hot/Cold DM (\\cite{sch92}; \\cite{dav92}; \\cite{tay92}; Klypin et al. 1993; Pogosyan \\& Starobinsky 1993) as a  DM model satisfying a wide body of observational data for clustering ( distributions of galaxies, galaxy clusters, quasars, Ly-$\\alpha$ systems, etc.) on typical scales $\\leq  l_{LS} \\sim 100-150 {h}^{-1}Mpc$ and with an Harrison-Zeldovich spectrum of the primordial cosmological perturbations on scales $>l_{LS}$. Because of the changes introduced into the power spectrum by neutrino clustering, for SMDM models it was shown the possibility to reconcile the evidence of high coherent velocity flows on large scales ($\\sim 50 {h}^{-1}Mpc$ ) with the moderate galactic pair velocities on megaparsec scales (\\cite{kly93}). For the cluster correlation function  ( \\cite{hol93}; \\cite{jin93}; \\cite{klr93}) the two-point function for SMDM models appears to be consistent with data  for $R \\magcir 20-30 h^{-1} Mpc$ ( on the contrary to what found for CDM ).  Analytical approximations for the present epoch transfer functions are given by Holtzmann (1989);  Pogosyan \\& Starobinsky (1995); Ma (1996); Eisenstein \\& Hu (1997). Early numerical simulations have considered $\\Omega_H=0.3$, we will use the notation $H$ for massive neutrinos, but this model does not produce Ly-$\\alpha$ systems as much as observed (\\cite{moa94}). For this reason a value of $\\Omega_H \\mincir 0.2$ results in a better fit (\\cite{kly95}; \\cite{maa94}). In comparison with SCDM or SHDM, MDM models are more complicated, we have the choice of two independent parameters : the ratio $\\Omega_H/\\Omega_C$ and $m_H$ ( if one allows for more than one species of massive neutrinos). The spectral index $n$ of the post-inflationary spectrum is taken $n=1$. It is possible to consider also the role of gravitational waves, which can change significally the normalization of the spectrum and the formation of structures (e.g.,  \\cite{maa96} ).  It must be stressed that SMDM models are in difficulty with the present upper limits  on the age of the Universe: if we assume $h=0.5$ and the age of globular clusters greater than $15 Gyr$, in the case of a flat FRW  model the needed value of $\\Omega_M $ is $\\leq 0.6 $. The other difficulties are connected with later galaxy and quasar formation (\\cite{pog95}; \\cite{cen94}). Another difficulty for SMDM models is that when the power spectra are normalized to the COBE 4-yr data the linear theory overpredicits the observed cluster abundances. The estimated uncertainties in the normalization and linear calculations can hardly fit the cluster number density within present data error bars. The difficulty can be reduced if one removes the constraint of a scale invariant spectra and introduces a small tilt ( $ n \\simeq 0.8 -0.9$, \\cite{maa96}).  In alternative to MDM models spatially flat low density models with a positive cosmological constant are considered to be a viable generalization of SCDM after COBE. These models are termed $\\Lambda$CDM and were considered even before the quadrupole detection (\\cite{pee84} and references therein). There are two main reason to consider $\\Lambda$CDM models as an alternative to MDM: the present lower limit on the age of the Universe and the baryon fraction in galaxy clusters.  Current uncertainties for the age of globular clusters give $t_0=15 \\pm 2 Gyr$ (\\cite{cha96}). If one takes a lower limits $t_0 \\geq 13 Gyr$ then in the $\\Omega_M=1$ case the maximum value for $h$ is $0.5$. This value is below the range allowed by recent HST measurement : $h=0.7\\pm.1$ (\\cite{fre94}; \\cite{rie95}). The introduction of a positive cosmological constant allows $h$ to be higher for a fixed age $t_0$ than in $\\Omega_M=1$. The other reason to consider $\\Lambda$CDM model is from X-ray observations of galaxy clusters. If clusters are a representative sample of the matter content in the Universe, then their baryons/total matter ratio should not be different from what expected from standard nucleosynthesis (\\cite{whb93}). For $h \\magcir 0.5 $ this is achieved if $\\Omega_M \\mincir  0.35 \\pm 0.2 $. Quite interestingly this range is close to that obtained from dynamical estimates. After COBE detection of the quadrupole  anisotropy $\\Lambda$CDM models have been further considered (Kofman, Gnedin \\& Bahcall 1993; Klypin, Primack \\& Holtzman 1996; Liddle et al. 1996a; Liddle et al. 1996b). They~can fit several constraints on large scales ($k \\mincir 1 {h}^{-1} {Mpc}^{-1} $, \\cite{lib96}), but small scales clustering it is excess (\\cite{kly96}) by a factor $\\approx 2-3$ when compared to estimates from galaxy catalogs,  so this is a difficulty of the model which can be solved with the introduction of an antibias of small scales. A possibility which seems unplausible (\\cite{pri96}).  These difficulties have suggested that consistency with present data can be achieved for standard DM models with the introduction of one extra parameter. Possible variants are two species of massive neutrinos ( Primack 1997 ), or a tilt of the primordial spectrum ( Cole et al. 1997 ).  The main aim of this paper is to consider an alternative possibility , that is MDM models with $n=1$ and a non-zero cosmological constant (MLM). These models have in fact the following free parameters : $\\Omega_{H}/\\Omega_M$, $\\Omega_\\Lambda$ and the number of species of massive neutrinos. If one considers one species of massive neutrinos, then MLM are on the same foot of the previous alternatives, the main advantage being that they retain the inflation paradigm and a scale-invariant initial spectrum.   In order to study these models we have numerically integrated forward in time the linearized Einstein-Boltzmann equations for the evolution of the metric and density perturbations. The final power spectra have been normalized according to the four year COBE data (\\cite{bun97}). We have then used linear perturbation theory to find the parameters for which MLM models are consistent with a variety of current observational constraints on large scales (power spectrum, cluster-cluster correlations, bulk velocities, cluster abundance). In order to make comparison with observations we will consider the following set of parameter space: $\\Omega_\\Lambda=0, 0.31, 0.45, 0.65, 0.74$; $\\Omega_H/\\Omega_M=0.1, 0.2, 0.3$; number of massive and massless neutrinos $\\beta_H=2, 4, 6$ and $\\beta_{\\nu}=0, 2, 4, 6$ ( here $\\beta$ is the sum of spin states ). In some cases we considered a number of neutrinos species greater than three, this possibility being of no physical relevance, we ran these models for studying  how different numbers  of massive and massless particles  would change the spectra evolution. Our final analysis will be restricted to one massive neutrino. The fifth parameter is $h$ ($h=0.4, 0.5, 0.63, 0.7$).  The outline of the paper is as follows: in Section 2 we will present our model description and equations describing the linear perturbations in MLM models. In Section 3 we will study the power spectra dependence on model parameters using COBE 4-year data normalization. In Section 4 some linear tests are given for MLM models, including our results for the cluster mass and temperature function. Our conclusions are given in Section 5. ",
        "conclusions": "We have discussed linear clustering evolution  for a set of spatially flat MDM models with a cosmological constant. We have not considered the role of gravitational waves or of a possible tilt in the primordial spectrum. In order to restrict the range of allowed values for the cosmological parameters, we have applied linear perturbation theory to compare the predictions of our MLM models with a set of linear data. The models considered had one species of massive neutrinos and a scale-invariant spectrum.  Linear calculations for the dimensionless power spectra $\\Delta^2(k)$ have been compared with the reconstructed real-space power spectrum of Peacock \\& Dodds (1994), Peacock (1997). Because we have considered linear spectra evolution a comparison of clustering is meaningful only at low $k$ ($ \\mincir 0.2h Mpc^{-1})$ A substantial agreement is obtained for those unbiased models with $\\Omega_H/\\Omega_M \\mincir 0.2$, and $\\Omega_M \\magcir 0.6$ ($ h = 0.5 $) or $ 0.3 \\mincir \\Omega_M \\mincir 0.5$ ($h=0.7$).  The computed linear power spectra have been filtered with a Gaussian window of radius $R_f=12h^{-1}Mpc$ to calculate the peculiar velocity field. A comparison of the rms bulk motion with POTENT data does not lead to strong constraints : $\\oml \\mincir 0.7$ and $\\Omega_H /\\Omega_M \\mincir 0.2$. The cluster-cluster correlation functions calculated for MLM spectra in the framework of Gaussian random density fluctuation field explain the observed positive correlations at $>50\\;h^{-1}Mpc$ and are within the limits of the error bars for that observed at large distances when $0.13\\le\\Omega_M h\\le 0.35$.  Of the considered tests the most important turns out to be that based on cluster abundances, for which the present X-ray cluster temperature function can put strong limits on the initial spectrum amplitude, or conversely on $\\sigma_8$. Because of the strong dependence of $N_{cl}(> kT)$ on temperature, the conversion of X-ray temperatures into masses is critical . From the discussion of Section 4.5 this conversion can be obtained with a dispersion of less than 10 \\%. We have then used the Press \\& Schechter (1974) formalism to compute the cluster number density for different MLM models. Consistency with present data  can be achieved for COBE normalized models only for the following range of values (at $1\\sigma$ level): $\\Omega_H /\\Omega_M \\le 0.2$, $ h =0.5(0.7)$ and $ 0.45(0.3) \\le \\Omega_M \\le 0.75(0.5)$. In correspondence with these limits, the allowed values of $\\sigma_8$ are : $\\sigma_8 \\simeq 0.65 \\pm 0.1$  ($h=0.5$) or $\\sigma_8\\simeq 0.8 \\pm 0.1$ ($h=0.7$). These limits suggest that MLM models require a moderate amount of bias to be consistent with clustering data. For these range of values consistency with clustering data yields $ 0.55(0.35) \\le \\Omega_M \\le 0.65(0.4)$ for $ h =0.5(0.7)$ and $\\Omega_H /\\Omega_M = 0.1$. Models with $\\Omega_H /\\Omega_M = 0.2$ are almost inconsistent.   One can also compare these constraints with those given by considering cluster evolution. The evolution of galaxy clusters has recently been used also for determining the cosmological parameters $\\sigma_8$ and $\\Omega_M$. Fan, Bahcall \\& Cen (1997) found $\\sigma_8=0.83\\pm 0.15$. This range of values is consistent with the $\\sigma_8$ required for MLM with $h=0.7$ (Fig. 15), and marginally consistent if $h=0.5$. At the $2\\sigma$ level none of the MLM models considered here is ruled out. Finally Henry (1997) has used recent data for the evolution of the  cluster X-ray temperature function to determine $\\Omega_M$ and $\\sigma_8$. For a flat cosmology he obtains $\\Omega_M=0.55 \\pm 0.17$ and $\\sigma_8=0.66^{+0.34}_{-0.17}$, in close agreement with our findings for MLM models.  The observed abundance of damped Lyman-$\\alpha$ systems has been used to obtain lower limits on the rms mass fluctuation $\\sigma(M)$ at $z=4$ using Press-Schechter theory. These limits yield the following constraints on $\\Omega_M$ for MLM models with $\\Omega_H /\\Omega_M \\le 0.2$ : $0.8-0.4 (0.3) \\mincir \\Omega_M $ for $h=0.5$ ($h=0.7)$. These values have been obtained assuming the fraction of neutral hydrogen to be unity. If this fraction is close to $10 \\%$, as suggested by numerical hydro simulations, then the constraints for MLM models become much more severe. In this case consistency with both damped Lyman-$\\alpha$ and cluster abundances can be achieved only for $h=0.7$ , $\\Omega_M \\simeq 0.4$ and $\\Omega_H /\\Omega_M =0.1$.   In order to further restrict our models we can also consider present observational constraints on cosmological parameters. The cosmological constant is constrained to be $\\oml \\le 0.66$ at the 95 \\% confidence limit by QSO lensing  (Kochanek 1996). A tighter restriction comes  from the recent work of Perlmutter et al. (1997) on SN Ia , who give $\\Omega_{\\Lambda} < 0.51$ still at 95 \\% c.l. .  Recent Hipparcos data ( Feast \\& Catchpole 1997, Reid 1997) have brought down the Cepheid scale distance, thus reducing the estimated Globular Cluster age to $t_0 > 12 Gyr$ and $h$ to  a mid term $h=.65$. These values still require a cosmological constant, but not as high as in $\\Lambda$CDM models. A value of $\\oml > 0.4 $ is needed to satisfy the new age constraint.  An upper limit on $\\Omega_M$ can be obtained from the estimated baryonic content of galaxy clusters. If clusters are massive enough to represent a fair sample of the total matter content in the universe, as numerical simulations confirm (\\cite{evr96}), then the baryon fraction $f_b=(0.06 \\pm 0.003) h^{-3/2}$  (Evrard 1997) should be close to its universal value.  Thus the standard Big Bang Nucleosynthesis value $\\Omega_b=(0.008-0.024)h^{-2}$ (\\cite{cop95}) can be used to infer $\\Omega_M=\\Omega_b /f_b = (0.25 \\pm 0.15) h^{-1/2}$. For $h=0.65$ one obtains $ \\Omega_M < 0.5$. For a flat model this limit already overlaps the lower limit from SN Ia. This is the main argument for a cosmological constant, possible counter arguments like a magnetic field pressure or density inhomogeneities, which can lead to underestimate the total cluster mass, are unlikely to push the limit up to $\\Omega_M=1$ ( Evrard 1997 and references cited therein ).  The region of parameter space which is allowed by these observational constraints for a flat model is then: $ 0.5 \\le h \\le 0.7$, $ 0.3 \\le \\Omega_M \\le 0.6$. These are also the limits that for MLM models the cosmological parameters must independently satisfy in order to achieve consistency with the set of linear clustering data previously analysed.  We think that this is a notable feature of MLM models  and one of the most important results of this paper. We then summarize our conclusions by saying that MLM models \\footnote{While this work was being completed, a paper appeared on babbage by Eisenstein \\& Hu (1997), who have considered power spectra for CDM and other DM models, including MLM.} appear to be a promising class of cosmological dark matter models. Our linear analysis shows that  consistency for the cosmological parameters is achieved over a wide range of observational data.   \\bigskip \\bigskip \\bigskip   {\\bf Acknowledgements}  T. Kahniashvili is grateful to ICTP and SISSA for financial support and B. Novosyadlyj also acknowledges financial support by  SISSA. T.K. and B.N. are grateful to SISSA for hospitality and the stimulating academic atmosphere which allowed this work to progress. RV thanks also S. Bonometto for helpful discussions."
    },
    "9804/astro-ph9804327_arXiv.txt": {
        "abstract": "As a class, LINERs and Low-Luminosity AGN tend to show little or no significant short-term variability (i.e., with time-scales less than a day). This is a marked break for the trend of increased variability in Seyfert 1 galaxies with decreased luminosity.  We propose that this difference is due to the lower accretion rate in LINERs and LLAGN which is probably causing the accretion flow to be advection-dominated.  This results in a larger characteristic size for the X-ray producing region than is the case in ``normal'' AGN.  Short-term variability may be caused by a localized instability or occultation events, but we note that such events would likely be accompanied by broad-band spectral changes. Since the ADAF is more compact in a Kerr metric, it is possible that the X-ray emission from ADAFs around rotating blackholes would be more variable than X-ray emission from ADAFs in a Schwarzchild metric. Similar variability arguments also apply to other wavelengths, and accordingly multiwavelength monitoring of LLAGN could serve to ``map'' the ADAF regions. ",
        "introduction": "From the time LINERs were originally identified as a class (\\cite{Heck80}) it has been debated whether or not they represent a low-luminosity version of AGN (see \\cite{fil96} for a review).  Recent optical work has shown that $\\sim 23\\%$ of LINERs have broad $H\\alpha$ lines (\\cite{Ho97b}), and X-ray observations of LINERs show that the 2-10 keV spectrum is dominated by a power-law with an energy index $\\alpha \\sim 0.8$ (\\cite{Ptak97} and references therein). These properties are highly suggestive that LINERs are indeed AGN in many cases.  Recent surveys have shown that LINERs and and other low-luminosity AGN (LLAGN; mostly Seyfert 2 galaxies) are very common in the nuclei of nearby galaxies, so it important to understand these objects.  Specifically, an important question is to what extent do the properties of ``normal'' AGN scale with luminosity or perhaps more precisely, what properties scale with accretion rate.  Modeling of low-accretion rate blackhole systems suggests that when the accretion rate drops below $\\sim 0.1-0.01$ in Eddington units (i.e., $\\dot{m}$ = $\\dot{M}/\\dot{M}_{Edd}$ , $\\dot{M}_{Edd}$ = $1.4 \\times 10^{18}\\frac{M}{M_{\\odot}} \\rm \\  g \\ s^{-1}$), the dominate mode of accretion will most likely be ``advection-dominated'' (\\cite{abram95}, \\cite{naryi95}). The radiative efficiency of an advection-dominated accretion flow (ADAF) is very low, unlike the optically-thick, geometrically-thin, disks thought to be present in normal AGN where the accretion efficiency is probably on the order of 0.1. In LLAGN and LINERs where the putative blackhole mass, $M_{BH}$, has been measured, $M_{BH}$ tends to be on the order of $10^{7-9}$ (c.f., $M_{BH} \\sim 3.5 \\times 10^7 \\rm M_{\\odot}$ in NGC 4258, \\cite{Green95}; $M_{BH} \\sim 10^9 \\rm M_{\\odot}$ in NGC 4594, \\cite{Kor96}). Luminosities on the order of $10^{40-41} \\rm \\ ergs \\ s^{-1}$ therefore imply luminosities and accretion rates on the order of $10^{-2}$ or less in Eddington units (assuming that on the order of 10\\% of the bolometric luminosity of these galaxies is radiated in X-rays), well within the ADAF parameter space.  Indeed, the presence of an ADAF in M81 (\\cite{Petre93}) and NGC 4258 (\\cite{L96}) has been suggested based in part on their low accretion rates (see also \\cite{mah97} and \\cite{yibou98}). However, we note that ADAF models for accretion rates below $10^{-3}$ predict that the X-ray band emission would be dominated by Bremsstrahlung emission from the $\\sim 10^{9-10}$ K electrons rather than by inverse-Compton scattering of synchrotron photons (c.f., \\cite{naryi95}, \\cite{yibou98}). With the exception of NGC 4258, thermal Bremsstrahlung fits to the 2-10 keV emission have resulted in temperatures on the order of $5-10 \\times 10^7 \\rm \\ K$, ruling out $10^9$ K emission.  Accordingly, the ADAF model predicts that the accretion rates in these galaxies is probably not much below $\\sim 10^{-3}$. As with normal AGN, the X-rays produced from the ADAF arise from closer in to the AGN and are attenuated less by absorption than emission at other wavelengths. In this letter we show that the X-ray temporal properties of LINERs and LLAGN are more consistent with an ADAF than with a geometrically-thin disk. ",
        "conclusions": "For some of the galaxies in this sample, the lack of variability is probably due to either the presence of multiple sources of the 2-10 keV emission or the fact that the hard emission is scattered at distances greater than a light-day from the nucleus.  The former case is most likely the situation for the starburst galaxies, where while some of the hard flux may be due to ``hidden'' micro-AGN, much of hard flux may to be due to multiple point-sources (supernovae and X-ray binaries) or hot gas (although as mentioned above M82 is similar to M81 in its variance). The latter case is most likely to be true for the Seyfert 2 galaxies NGC 3147, NGC 4258, and M51, as suggested in \\cite{Ptak96}.  However, note that in the case of NGC 4258, the observed column density is on the order of $10^{23} \\rm \\ cm^{-2}$ which is consistent with the column densities in Seyfert 2 galaxies observed by \\ginga (\\cite{Awaki92}) and \\asca (\\cite{turner97}).  While it is possible that the hard X-ray flux from NGC 4258 is scattered around material with a column on the order of $10^{24} \\rm \\ cm^{-2}$ or more, it is most likely that the 2-10 keV flux we are observing is the direct nuclear continuum (note that this argument is strengthened by the classification of NGC 4258 as a Seyfert 1.9 in \\cite{Ho97a}).  None of the above considerations are likely to be true for the brighter LINERs in this sample, M81, NGC 3998, and NGC 4579 since each of these has been observed to exhibit broad $H\\alpha$ emission (\\cite{Ho97b}), making them ``Type-I'' LINERs, analogues to the higher-luminosity Seyfert 1 galaxies.  In the cases of M81 (\\cite{ishi96}), NGC 4579 (\\cite{S96}) and, interestingly, the ``transition'' starburst-LINER galaxy NGC 3628 (\\cite{d95}; \\cite{y95}), significant variability has been observed between the \\rosat and \\asca observations, indicating that the nuclear sources are dominating the emission.  It is therefore likely that in most of these galaxies the dominant mode of accretion is fundamentally different from that in Seyfert galaxies, since it is accretion that is driving the X-ray luminosity and, by extension, the X-ray variability.  A fundamental difference between an ADAF and an optically-thick accretion disk is that in an ADAF it is the flow itself that is producing the X-rays.  The $\\alpha$-disk solution predicts a temperature of less than $10^{6}$ for $M_{BH} > 10^{7} \\ \\rm M_{\\odot}$ (\\cite{ss73}; see \\cite{frank92} for a review).  In this case (i.e., ``normal'' Seyfert galaxies), the X-ray continuum is most likely produced by the inverse-Compton scattering of UV photons from the ``cold'' accretion disk by energetic electrons, possibly in a ``corona'' above the disk (c.f., \\cite{HM91}). In an ADAF, the X-rays are produced by either the Comptonization of synchrotron radiation by the electrons in the flow or by Bremsstrahlung emission from the electron themselves.  In either case, the essence of the ADAF solution is that the emission mechanism is inefficient (on accretion time-scales) and a substantial volume contributes to the X-ray emission.  Since the ADAF is likely to be nearly spherical, most of the X-ray emission originates in a volume that is probably a spherical annulus extending from $r \\sim 5-10 \\ R_{Schw}$ in the case of a stationary blackhole or $r \\sim 3-7$ in the case of a maximally-rotating blackhole.  If variability is due to a change in $\\dot{m}$, then the {\\it minimum} time for the ADAF to respond is on the order of $\\pi\\bar{r}\\gamma_g/c$, where $\\bar{r} = \\frac{\\int L_X(r)rdr}{\\int L_X(r)dr}$, $L_X(r)$ is the luminosity of the ADAF in the 2-10 keV bandpass at $r$ and $\\gamma_g$ is the gravitational time dilation ($\\gamma_g \\sim [1-R_{Schw}/r]^{-1/2}$). $\\bar{r}/c$ is probably on the order of 2 and 6 $R_{Schw}/c$ (i.e., near the inner-most stable orbits), corresponding to $\\sim 4$ and $11\\frac{M_{BH}}{3.5 \\times 10^{7} \\rm \\ M_{\\odot}}$ ks light travel time-scales. More realistically, the time required for the ADAF to respond to a change in $\\dot{m}$ will depend on the sound speed, which approaches $0.25c$ (stationary blackhole) or $0.5c$ (rotating blackhole) at small radii (\\cite{JK97}), resulting in a time-scale on the order of $\\sim 6-20$ ks.  It is evident that a Kerr metric should result in a more variable ADAF, while {\\it any} ADAF onto a blackhole at least as massive as $\\sim 10^7 \\rm \\ M_{\\odot}$ should not be variable on time-scales less than several to tens of ks. The time-scales probed by the {\\it ASCA} observation are on this order.  This analysis is based on the assumption that the any variability would be caused by rapid changes in $\\dot{m}$ and that the ADAF will remain stable under these changes. However, detailed modeling by \\cite{Man96} and \\cite{Tak97} showed that instabilities produced by rapid changes in $\\dot{m}$ are likely to produce ``shots'' in the X-ray luminosity of an ADAF with time-scales on the order of several hundred $R_{Schw}/c$ (as observed in the low state of Cygnus X-1), suggesting that the time-scales assumed above are very conservative. It may be simply that the accretion rate itself is more steady in LLAGN than it is in Seyfert galaxies.  Note that this also would then indicate a break with Seyferts, assuming that changes in accretion rate are ultimately responsible for variability in Seyferts.  If occultation events are the cause of variability in Seyferts, then the accretion flows in LLAGN lack material subtending a sufficient solid angle to produce the same type of variability, again suggestive of a larger region responsible for the production of X-rays.  Note also that if variability on short time-scales is observed in a LLAGN believed to have an ADAF is due to either an occultation event or an instability (where the X-ray flux would be dominated by shocks rather than the steady ADAF flow emission), then detectable spectral variability is likely to accompany the event.  Note that many of the arguments in this letter apply equally well to other wavelengths, although with the exception of radio, the nuclear component of the emission from a typical galaxy is absorbed, difficult to segregate from extra-nuclear emission, or both.   However, since ADAF models predict that the contribution of synchrotron cooling (dominating the radio) and Compton and bremsstrahlung cooling (dominating in X-rays) to the ADAF luminosities varies only slightly as a function of radius, monitoring of the radio and X-ray luminosity of an ADAF should be further test of ADAF models.  M81 was observed to vary in the radio on time-scales of weeks (\\cite{biet98}, \\cite{Ho98}) and large-scale ($\\Delta I/I \\sim 1.7$) X-ray variability with a time-scale of months and 20\\% variability with time-scale of $\\sim 1$ day was observed by \\asca (\\cite{ishi96}, \\cite{S96}). Accordingly, the ADAF model predicts not only a specific broadband spectral shape but also that the radio and X-ray flux from galaxies such as M81 would be correlated over long periods of time (i.e., on time-scales sufficient for the ADAF to be in equilibrium).  Finally, note that as the accretion rate increases, the transition radius where the accretion flow changes from an ADAF to a thin-disk flow would decrease (see \\cite{Esin97}), possibly resulting in short-term variability as observed in Seyfert 1 galaxies.  In the case of hard X-ray emission, this short-term variability should be accompanied by the spectral features associated with $\\alpha$-disks (i.e., Fe-K emission and the Compton reflection ``hump'') which should not be present in AGN dominated by ADAFs."
    },
    "9804/astro-ph9804132.txt": {
        "abstract": "Three-dimensional hydrodynamical simulations are presented for the direct head-on or off-center collision of two neutron stars, employing a basically Newtonian PPM code but including the emission of gravitational waves and their back-reaction on the hydrodynamical flow. A physical nuclear equation of state is used that allows us to follow the thermodynamical evolution of the stellar matter and to compute the emission of neutrinos. Predicted gravitational wave signals, luminosities and waveforms, are presented. The models are evaluated for their implications for gamma-ray burst scenarios. We find an extremely luminous outburst of neutrinos with a peak luminosity of more than $4\\cdot 10^{54}\\,{\\rm erg/s}$ for several milliseconds. This leads to an efficiency of about 1\\% for the annihilation of neutrinos with antineutrinos, corresponding to an average energy deposition rate of more than $10^{52}\\,{\\rm erg/s}$ and a total energy of about $10^{50}\\,{\\rm erg}$ deposited in electron-positron pairs around the collision site within 10$\\,$ms. Although these numbers seem very favorable for gamma-ray burst scenarios, the pollution of the $e^\\pm$ pair-plasma cloud with nearly $10^{-1}\\,M_{\\odot}$ of dynamically ejected baryons is 5 orders of magnitude too large. Therefore the formation of a relativistically expanding fireball that leads to a gamma-ray burst powered by neutrino emission from colliding neutron stars is definitely ruled out. ",
        "introduction": "The {\\it merging} of two neutron stars that make up a binary system is of interest both because it is a powerful source of gravitational waves and because it might be the central engine of gamma-ray bursts. These mergings have been studied intensely during the past few years, since the occurrence rate is high enough (e.g.~Narayan et al.~1991, Phinney 1991, Tutukov et al.~1992, Tutukov \\& Yungelson 1993, Lipunov et al.~1995, van den Heuvel \\& Lorimer 1996, Lipunov et al.~1997, Prokhorov et al.~1997, Bethe \\& Brown 1998) that one might be able to measure the observable consequences. However, direct {\\it collisions} of two neutron stars are much rarer [the rate is $\\ll 10^{-10}$ per year per galaxy (Centrella \\& McMillan 1993) as compared to the binary merger rate of  $10^{-6}$ to $10^{-5}$ per year per galaxy] and thus have not received much attention.  Two previously published sets of hydrodynamical models of colliding neutron stars were computed by Centrella \\& McMillan (1993) and by Rasio \\& Shapiro (1992). Both used an SPH code to simulate the dynamics of two polytropic neutron stars of equal masses with adiabatic index $\\gamma=2$ and started at an initial distance of about four neutron star radii on parabolic orbits. Using a relatively small number of particles (2048), Centrella \\& McMillan (1993) were able to produce a catalog of gravitational wave luminosities and gravitational waveforms for different impact parameters. Rasio \\& Shapiro (1992) restricted themselves to fewer simulations with higher resolution (16000 particles) of neutron star coalescence as well as head-on collisions. They found that up to 5\\% of the total mass can escape from the systems and that strong shocks occur. Their results for the gravitational waveforms and luminosities reveal a smaller first maximum (caused by the initial free fall of the two stars), followed by the main peak associated with the rapid deceleration of the colliding matter during the propagation of the recoil shocks.  Neutron star collisions have repeatedly been suggested in the literature as possible sources of gamma-ray bursts (e.g., Katz \\& Canel 1996, Dokuchaev\\& Eroshenko 1996; also Dokuchaev et al.~1998), powered either by neutrino-antineutrino annihilation which produces an $e^\\pm$ pair-\\-pho\\-ton fireball, or by highly relativistic shocks which are formed during the collision (or coalescence) and eject matter at relativistic velocities (Shaviv \\& Dar 1995). Katz \\& Canel (1996), in particular, developed the idea that long gamma-ray bursts might be explained by accretion induced collapse of bare degenerate white dwarfs, while short bursts might originate from neutron star {\\it collisions}. They estimated the post-collision temperature to be $kT \\approx 100$~MeV, and expected a neutrino pulse with a width of 2.5$\\,$ms, and a total number of escaping neutrinos of $10^{58}$ (refering to Dar et al.~1992). With a mean neutrino energy of about $6\\,$MeV this corresponds to a total energy of $10^{53}$~erg emitted in neutrinos. %%%%%%$10^{58}$ * 6MeV * 1.6E-6erg/MeV = $10^{53}$ erg(!) They sketched the picture that shortly after the collision, dense lumps of hot matter expand and become optically thin to neutrinos which are released in a powerful outburst that is luminous enough to produce the desired gamma-ray burst energy by neutrino-antineutrino annihilation if one assumes an efficiency of 1\\%. In order to obtain the observed gamma-ray burst rates they had to invoke a total of $10^9$ hypothetical clusters within $z\\approx 1$ with $10^8$~neutron stars each. They stated that ``we cannot exclude the possibility that such clusters are commonly found at the centers of galaxies''.  Neutron star collisions might also be considered as interesting, because they could be viewed at as more violent variants of the merging process with a more extreme dynamical evolution. This might support the formation of relativistic shocks and possibly lead to a larger burst of gravitational waves and neutrinos. More violence of the merging process can be expected from general relativistic effects which were so far neglected in the large majority of simulations (see, however, Wilson et al.~1996 for an interesting step towards fully general relativistic modeling). Oohara \\& Nakamura (1997), using a grid-based TVD scheme, performed preliminary simulations to compare the coalescence with a Newtonian potential to the case with post-Newtonian physics. Their neutron stars were modeled as two polytropes with $\\gamma=2$, $M=0.62\\,M_\\odot$, and $R=15\\,$km. The principal difference between the two cases was that the post-Newtonian merging indeed turned out to be more violent: the impact is more central and shocks develop, because ``general relativity effectively increases the gravitational force''. Although the strong shock itself has little immediate effect on the gravitational waveform, the dynamics can be changed because of the higher temperatures and densities. Wex (1995) and Ogawaguchi \\& Kojima (1996) showed that both spin-orbit and spin-spin interactions appear formally to be {\\it first} order post-Newtonian corrections, just as gravitational potential corrections are (gravitational waves are 2.5$\\,$PN), but the inferred magnitude of these corrections for known compact binary systems is actually smaller.    %--------------------------- Table 1 ------------------------------------- \\begin{table*} \\caption[]{ Characterizing parameters and some computed quantities for all models. $N$ is the number of grid zones per dimension, $L$ the size of the largest grid, $l$ the size of the smallest zone, $M_{\\rho<11}$ the gas mass with density below $10^{11}$~g/cm$^3$ at the end of the simulation, and $M_{\\rm d}$ the mass with specific angular momentum larger than $j^*\\equiv v_{\\rm Kepler}(3R_{\\rm s})3R_{\\rm s}$ where $R_{\\rm s}$ is the event horizon of a $\\sim 3\\,M_{\\odot}$ black hole probably forming from the colliding neutron stars. $M_{\\rm g}$ denotes the mass swept off the grid, $M_{\\rm u}$ is the mass leaving the grid unbound, and $T_{\\rm ex}$ the maximum temperature (in energy units) reached on the grid during the simulation of a model. }  \\begin{flushleft} \\tabcolsep=3mm \\begin{tabular}{lllcccccccc} \\hline\\\\[-3mm] Model & impact & spin & $N$ & $L$ & $l$ & $M_{\\rho<11}$ & $M_{\\rm d}$ & $M_{\\rm g}$ & $M_{\\rm u}$ & $T_{\\rm ex}$ \\\\ & direction & & & km & km & {\\scriptsize$10^{-2}M_\\odot$} & {\\scriptsize$10^{-2}M_\\odot$} & {\\scriptsize$10^{-2}M_\\odot$} & {\\scriptsize$10^{-2}M_\\odot$} & {\\scriptsize MeV} \\\\[0.3ex] \\hline\\\\[-3mm] h & head-on       & none & ~32 & 328 & 1.28 & 4.0 & 0.0 & 5.2 & 1.5 & 96.                     \\\\   % Grbc H & head-on       & none & ~64 & 328 & 0.64 & 4.1 & 0.0 & 6.6 & 1.5 & 96.                     \\\\   % GrbC ${\\cal H}$& head-on  & none & 128 & 328 & 0.32 & --- & --- & --- & --- & --- \\\\[0.7ex]  % GrBC o & off-center    & none & ~32 & 328 & 1.28 & 9.0 & 0.03 & 1.5 & 0.16 & 57.                     \\\\   % Grbo O & off-center    & none & ~64 & 328 & 0.64 & 11.0& 0.03 & 1.8 & 0.19 & 58. \\\\[0.7ex]  % GrbO \\hline \\end{tabular} \\end{flushleft}  \\label{tab:models1} \\end{table*} %--------------------------- Table 1 end ---------------------------------   %--------------------------- Table 2 ------------------------------------- \\begin{table*} \\caption[]{ Gravitational-wave and neutrino emission properties for all models. $\\widehat{\\cal L}$ is the maximum gravitational-wave luminosity, ${\\cal E}$ the total energy emitted in gravitational waves, $r\\widehat{h}$ the maximum amplitude of the gravitational waves as observed from a distance $r$, $L_{\\nu_e}$ the stationary value of the electron neutrino luminosity which is reached at about 6--10~ms after the start of the simulations, $L_{\\bar{\\nu}_e}$ the corresponding electron antineutrino luminosity, and $L_{\\nu_x}$ the luminosity of each individual species of $\\nu_x$ ($ = \\nu_{\\mu},\\,\\bar\\nu_{\\mu},\\,\\nu_{\\tau},\\,\\bar\\nu_{\\tau}$). $L_\\nu$ represents the total neutrino luminosity at the end of the simulation, $\\langle\\epsilon_{\\nu_e}\\rangle$, $\\langle\\epsilon_{\\bar{\\nu}_e}\\rangle$ and $\\langle\\epsilon_{\\nu_x}\\rangle$ are the mean energies of the different neutrino and antineutrino flavors. $\\dot{E}_{\\nu\\bar{\\nu}}$ denotes the integral rate of energy deposition by neutrino-antineutrino annihilation, averaged over the simulation time of 10$\\,$ms. }  \\begin{flushleft} \\tabcolsep=2mm \\begin{tabular}{llccccccccccc} \\hline\\\\[-3mm] Model & impact & $\\widehat{\\cal L}$ & ${\\cal E}$ & $r\\widehat{h}$ & $L_{\\nu_e}$ & $L_{\\bar{\\nu}_e}$ & $L_{\\nu_x}$ & $L_\\nu$ & $\\langle\\epsilon_{\\nu_e}\\rangle$ & $\\langle\\epsilon_{\\bar{\\nu}_e}\\rangle$ & $\\langle\\epsilon_{\\nu_x}\\rangle$ & $\\dot{E}_{\\nu\\bar{\\nu}}$  \\\\ & direction & {\\scriptsize$10^{55}\\frac{\\rm erg}{\\rm s}$} & {\\scriptsize$10^{52}$ erg} & {\\scriptsize$10^4$cm} & {\\scriptsize$10^{53}\\frac{\\rm erg}{\\rm s}$} & {\\scriptsize$10^{53}\\frac{\\rm erg}{\\rm s}$} & {\\scriptsize$10^{53}\\frac{\\rm erg}{\\rm s}$} & {\\scriptsize$10^{53}\\frac{\\rm erg}{\\rm s}$} & {\\scriptsize MeV} & {\\scriptsize MeV} & {\\scriptsize MeV} & {\\scriptsize$10^{50}\\frac{\\rm erg}{\\rm s}$} \\\\[0.3ex] \\hline\\\\[-3mm] h & head-on & 4.2  & 0.42 & 6.3 & 2.4 & 5.0  & 1.4  & 13. & 14. & 20. & 26. & --- \\\\   % Grbc H & head-on & 3.61 & 0.39 & 6.2 & 2.0 & 4.0  & 1.0  & 10. & 13. & 20. & 25. & 100 \\\\   % GrbC ${\\cal H}$& head-on & 3.55 & ---  & 6.2 & --- & ---  & ---  & --- & --- & --- & --- & --- \\\\[0.7ex]  % GrBC o & off-center & 6.6 & 4.1 & 12.3 & 1.3  & 2.8  & 0.8  & 7.3 & 11. & 18. & 25. & --- \\\\   % Grbo O & off-center & 6.1 & 3.6 & 11.9 & 1.2  & 3.1  & 0.8  & 7.5 & 11. & 18. & 25. & --- \\\\[0.7ex]  % GrbO \\hline \\end{tabular} \\end{flushleft}  \\label{tab:models2} \\end{table*} %--------------------------- Table 2 end ---------------------------------   Our project of simulating neutron star collisions with a Newtonian PPM code was motivated by the aspects described in the last two paragraphs. On the one hand, we intended to put potential gamma-ray burst scenarios to a test, on the other hand we wanted to study a situation that mimics the {\\it merging} of two neutron stars with extreme violence and maximal parameters like pre-merging kinetic energy and angular momentum in the system. Thus we hoped not only to obtain an upper bound on the gravitational wave emission to be expected from the merging of two neutron stars, even with general relativistic effects included. We also wanted to see whether the most extreme conditions during the collision of the two stars lead to sufficiently large neutrino emission to explain the gamma-ray burst energetics by the annihilation of neutrinos and antineutrinos emitted during the dynamical event. The latter seems impossible in case of the final stages of the coalescence of binary neutron stars because the prompt neutrino burst, although very luminous, fails by several orders of magnitude to produce about $10^{51}\\,{\\rm erg}$ of gamma-rays within the short time of only a few milliseconds that it takes the two neutron stars to merge into one massive body that is most likely going to collapse into a black hole on a dynamical timescale (Ruffert et al.~1996).  %Prokhorov M.E., Lipunov V.M., Postnov K.A., 1997, astro-ph/9704039:\\\\ %summarizing overview of binary NS merging rates: observational %($10^{-6}$) and theoretical ($10^{-4}$--$10^{-5}$).  The paper is organized as follows. In Sect.~\\ref{sec:numer} the basic aspects and new elements of our computational method (in extension of Ruffert et al.~1996) and the chosen initial conditions for our simulations are given. Our results are presented in the following sections, where we describe the hydrodynamical and thermodynamical evolution (Sect.~\\ref{sec:hydevol}), the gravitational wave emission (Sect.~\\ref{sec:gwem}), and the neutrino production (Sect.~\\ref{sec:nuem}) in the colliding stars together with the evaluation of our models for the efficiency of neutrino-antineutrino annihilation (Sect.~\\ref{sec:annihi}). Section~\\ref{sec:end} contains a summary and conclusions.    \\begin{figure*} \\tabcolsep=0.0mm \\begin{tabular}{cc} \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H1dens.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf a}}   & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H2dens.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf b}}  \\\\[-2ex] \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H1temp.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf c}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H2temp.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf d}}  \\\\[-4ex] \\end{tabular} \\caption[]{Contour plots of Model~${\\cal H}$ (left panels) and~H (right panels) showing cuts in a plane containing the $x$-axis which is the symmetry axis of the initial model. The displayed physical quantities are the density together with the velocity field (panels a and b) and the temperature (panels c and d). The density is measured in ${\\rm g\\,cm}^{-3}$, the temperature in MeV. The density contours are spaced logarithmically with intervals of 0.5~dex, while the temperature contours are linearly spaced, starting with 1$\\,$MeV, 3$\\,$MeV, 5$\\,$MeV, and then continuing with an increment of 5~MeV. The bold contours are labeled with their corresponding values ($10^{10}$, $10^{12}$, and $10^{14}\\,{\\rm g\\,cm}^{-3}$, and 10, 30, and 50$\\,$MeV, respectively). In the box in the upper right corner of each panel, the velocity vectors and the time elapsed since the beginning of the simulation are given. The mirror symmetry relative to the plane $x = 0$ and the rotational symmetry around the $x$-axis are broken during the evolution (panels b and d) because of an instability of the contact layer of the two neutron stars against shear motions by which numerical fluctuations (panel c) are amplified. } \\label{fig:Hcont1} \\end{figure*}  \\begin{figure*} \\tabcolsep=0.0mm \\begin{tabular}{cc} \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H1Ye.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf a}}   & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H2Ye.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf b}}  \\\\[-2ex] \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H1ent.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf c}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./H2ent.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf d}}  \\\\[-4ex] \\end{tabular} \\caption[]{Contour plots of Model~${\\cal H}$ (left panels) and~H (right panels) showing cuts in a plane containing the $x$-axis which is the symmetry axis of the initial model. The displayed physical quantities are the electron fraction $Y_e$ (panels a and b) and the entropy (panels c and d), the latter measured in units of Boltzmann's constant $k$ per nucleon. The contours of the electron fraction are linearly spaced with intervals of 0.02 below 0.1 and with intervals of 0.05 above, the entropy contours are given in steps of $1\\,k$/nucleon from 1 to 6, in steps of $2\\,k$/nucleon between 6 and 16, and then for values of 20, 25, 30, and 40$\\,k$/nucleon. The bold contours are labeled with their corresponding values (0.02, 0.06, 0.10, 0.20, 0.30, and 0.40 for $Y_e$, and 6, 10, and 20 for the entropy). Maximum values of $Y_e$ are above 0.4, of the entropy near $30\\,k$/nucleon. In the box in the upper right corner of each panel, the time elapsed since the beginning of the simulation is given. } \\label{fig:Hcont2} \\end{figure*}  \\begin{figure*} \\tabcolsep=0.0mm \\begin{tabular}{cc} \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O1dens.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf a}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O2dens.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf b}} \\\\[-2ex] \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O1temp.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf c}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O2temp.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf d}} \\\\[-4ex] \\end{tabular} \\caption[]{Contour plots of Model~O showing cuts in the orbital plane for the density together with the velocity field (panels a and b) and for the temperature (panels c and d). The density is measured in ${\\rm g\\,cm}^{-3}$, the temperature in MeV. The density contours are spaced logarithmically with intervals of 0.5~dex, while the temperature contours are linearly spaced, starting with 1$\\,$MeV, 3$\\,$MeV, 5$\\,$MeV, and then continuing with an increment of 5~MeV. The bold contours are labeled with their corresponding values ($10^{10}$, $10^{12}$, and $10^{14}\\,{\\rm g\\,cm}^{-3}$, and 10, 30, and 50$\\,$MeV, respectively). In the box in the upper right corner of each panel, the velocity vectors and the time elapsed since the beginning of the simulation are given. } \\label{fig:Ocont1} \\end{figure*}  \\begin{figure*} \\tabcolsep=0.0mm \\begin{tabular}{cc} \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O1Ye.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf a}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O2Ye.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf b}} \\\\[-2ex] \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O1ent.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf c}}  & \\epsfxsize=8.8cm \\epsfclipon \\epsffile{./O2ent.ps} \\put(-0.92,6.75){{\\Large \\bf \\sf d}} \\\\[-4ex] \\end{tabular} \\caption[]{Contour plots of Model~O showing cuts in the orbital plane for the electron fraction $Y_e$ (panels a and b) and the entropy (panels c and d), the latter quantity measured in units of Boltzmann's constant $k$ per nucleon. The contours are linearly spaced with intervals of 0.02 for the electron fraction and $1\\,k$/nucleon for the entropy. The bold contours are labeled with their corresponding values (0.02, 0.06, 0.10, 0.20 for $Y_e$ and 5 and 10 for the entropy). Maximum values of $Y_e$ are around 0.16, of the entropy about $10\\,k$/nucleon. In the box in the upper right corner of each panel, the time elapsed since the beginning of the simulation is given. } \\label{fig:Ocont2} \\end{figure*} ",
        "conclusions": "  \\noindent 1.) {\\it Dynamical evolution:} During the head-on collision of the two neutron stars, the collision remnant is heated by recoil shocks and the initial kinetic energy is efficiently shock-dissipated to thermal energy within a few violent pulsations on a timescale of about 3--4$\\,$ms, after which the collision remnant has an essentially spherical mass distribution. The neutron star matter is transiently heated to peak temperatures close to $100\\,$MeV. Average temperatures are around 40--50$\\,$MeV, corresponding to entropies between 2$\\,k$/nucleon and 10$\\,k$/nucleon. The violent crash of the stars into each other leads to the dynamical ejection of about 0.5\\% of the system mass. In contrast, the off-center collision is much ``milder'' and the ejected mass is about a factor of 10 smaller, although the collision remnant retains a large fraction (about 30\\%) of the initial kinetic energy as rotational energy even after 10$\\,$ms at the end of our simulations. Only for extreme assumptions about the nuclear EOS will the $\\ga 3\\,M_{\\odot}$ remnant of the head-on collision not form a black hole on a dynamical timescale. We also believe that the remnant of the off-center collision will probably not be able to escape the collapse to a black hole within a few milliseconds. A mass of $3\\,M_{\\odot}$ is distributed within a radius of 30$\\,$km and $2.8\\,M_{\\odot}$ are within 20$\\,$km (Fig.~\\ref{fig:rotparO}) which is only about twice the Schwarzschild radius of a $3\\,M_{\\odot}$ black hole. Thermal pressure can only insignificantly raise the maximum stable mass of neutron stars (Goussard et al.~1997) and rotation is able to increase the stable mass limit by only $\\la 20\\%$ (Friedman et al.~1986). The interior $\\sim 2.8\\,M_{\\odot}$ of the remnant rotate nearly uniformly and the relativistic rotation parameter of this mass is only $a(r=20\\,{\\rm km})=\\eck{Jc/(GM^2)}_{20\\,{\\rm km}} \\approx 0.3$ (cf.~Fig.~\\ref{fig:rotparO}), which is smaller than that of the maximally rotating, maximum-mass models constructed by Friedman et al.~(1986). For all but one extreme EOS tested by Friedman et al.~(1986) such a configuration is unstable. Therefore we conclude that it is very likely that gravitational instability will set in as soon as the two stars have merged into one massive body within $t\\ga 3$--$4\\,$ms. General relativistic effects will certainly influence the transformation of the infall orbit of the neutron stars into a bound one; this needs to be studied with relativistic simulations.  \\noindent 2.) {\\it Gravitational waves and neutrinos:} The gravitational-wave signal will certainly depend on general relativistic effects and our basically Newtonian models have only a limited ability to make predictions of a possibly measurable pulse. Moreover, the duration of the gravitational-wave and neutrino emission from the hot collision remnant will depend on the timescale of the delay until black hole formation. Our simulations yield a maximum amplitude of the gravitational waveform that is $h_{\\rm max}\\approx 2$--$4\\times 10^{-23}$ for neutron star collisions happening at a distance of $1\\,$Gpc. The off-center collision is the stronger gravitational-wave source due to the larger quadrupole moment of the rotating system and the longer duration of the emission which might last for 10--20 wave periods in the 1000--2000$\\,$Hz range. The gravitational-wave strain would be close to the lower sensitivity limit of the new generation of gravitational-wave interferometers which are currently under construction and will start operation within the next years. Of course, neutron star collisions are very rare and very short events and therefore the chance to catch a signal is rather small. The head-on collision is the more powerful neutrino source of the two investigated cases and emits an energy about 3 times larger in neutrinos than in gravitational waves. The peak neutrino luminosity reaches $4\\times 10^{54}\\,{\\rm erg\\,s^{-1}}$ and the total energy radiated in neutrinos within a few milliseconds is around $10^{52}\\,$erg.  \\noindent 3.) {\\it Gamma-ray bursts:} Because of the larger energy output in neutrinos and the very high neutrino luminosity ($L_{\\nu} \\ga 10^{54}\\,{\\rm erg\\,s}^{-1}$) as well as high mean energies of the emitted neutrinos ($\\ave{\\epsilon_{\\nu}}\\la 40\\,$MeV), the head-on collision provides more favorable conditions for producing gamma-ray bursts from $e^+e^-$-photon fireballs created by $\\nu\\bar\\nu$ annihilation. We calculate a conversion efficiency of neutrino energy into $e^+e^-$-pairs of about 1\\% and find an integral value for the energy deposited in the vicinity of the collision remnant of $10^{50}\\,$erg within only 10$\\,$ms. However, most of this energy is deposited in the immediate neighbourhood of the neutrinospheres where the density is still higher than about $10^{10}\\,{\\rm g\\,cm}^{-3}$. Therefore the baryon loading of the $e^+e^-$-photon fireball is at least 5 orders of magnitude too high and instead of having values above 100 the relativistic Lorentz factor is estimated to be only around $10^{-3}$. Dynamically ejected material together with a flow of baryonic matter driven by neutrino-energy transfer to the surface layers of the collision remnant are therefore a harmful poisonous combination which prevents relativistic expansion of the pair-plasma fireball even though the latter seems to obtain an interesting amount of energy from $\\nu\\bar\\nu$-annihilation.  Strong shock heating of the neutron stars during their violent collision (or during their merging as suggested by post-Newtonian simulations, see Oohara \\& Nakamura 1997) can indeed raise the neutrino luminosities significantly and can thus enhance the energy deposition by $\\nu\\bar\\nu$ annihilation. The associated dynamical ejection of gas and the neutrino-driven wind caused by the intense neutrino fluxes, however, impede the emergence of gamma-rays from this scenario. Fireballs powered by neutrinos from colliding neutron stars should therefore be ruled out as possible sources of cosmological gamma-ray bursts. Also, we do not see the formation of highly relativistic shocks during the collisions by which a fraction of 0.1--1\\% of the kinetic energy at impact ($\\ga 10^{53}\\,$erg) might be hydrodynamically focused into a small amount of mass ($\\sim 10^{-6}$ to $10^{-5}\\,M_{\\odot}$). The expansion velocity of the collision shocks is at most a few tenths of the speed of light and the masses ejected in the presented models are more than 3 orders of magnitude too large to allow for ultrarelativistic motion.    %Movies in mpeg format of the dynamical evolution of all models are %available in the WWW at %{\\tt http:\\nix//www.mpa-garching.mpg.de\\nix/\\lower0.7ex\\hbox{$\\!$\\~~$\\!$}mor\\nix/nsgrb.html}"
    },
    "9804/astro-ph9804333_arXiv.txt": {
        "abstract": "We present the method and first results of a survey for high-redshift ($z>3$) radio-loud quasars, which is based on optical identifications of 2902 flat-spectrum radio sources with $\\rm S_{5GHz}\\ge 25$mJy. The radio sample was defined over a 1600 sq degree region using the 5GHz Green Bank survey and the 1.4GHz VLA FIRST survey. 560 sources were identified to a limit of E=19.5 on APM scans of POSS-I plates and 337 of these optical counterparts are unresolved. From these a complete sample of 73 sources for spectroscopic follow up was defined based on criteria of red ($\\rm O-E \\ge 1.2$) optical colour.  We have obtained spectra for 36 of these and an additional 14 had redshifts in the literature, thus 70\\% of the spectroscopic sample is completed.  Six objects in the sample were found to be radio-loud quasars with $z>3$ of which two were previously known.  The efficiency of the spectroscopic phase of the survey is therefore about 1 in 9, whereas without the colour selection criterion the efficiency would have been 1 in $\\sim 40$.  The six $z>3$ quasars were found in an effective area of 1100 square degrees, implying a surface density of one flat-spectrum $z>3$ radio-loud quasar per 190 square degrees to limits of E=19.5 and $\\rm S_{5GHz}\\ge 25mJy$.  This survey also produced the first known radio-loud BAL quasar, 1556+3517 with $z=1.48$, which has been reported in an earlier paper (Becker et al 1997). This object has a redder optical colour ($\\rm O-E=2.56$) than all the $z>3$ quasars found in this survey to date. In addition we have obtained spectra of 22 GB/FIRST sources which are not part of the complete sample.  We give positions, E (red) magnitudes, $\\rm O-E$ colours, radio fluxes, radio spectral index and redshifts where possible for objects for which we have obtained spectra. We give spectra and finding charts for the $z>3$ quasars. ",
        "introduction": "We have begun a search for high-redshift quasars selected using the VLA FIRST survey (Becker, White \\& Helfand 1995). The aim of this work is to extend recent searches for radio-loud high-redshift quasars (e.g. Hook et al. 1995, 1996; Hook \\& McMahon 1998; Shaver et al. 1996) fainter in radio flux density by a factor of $\\sim 10$, and, ultimately, to use the new sample to constrain the faint end of the radio quasar luminosity function at high redshift. In addition, the sample can be used to provide an unbiased sample of damped absorption systems that can be used to study galaxy evolution at high redshift.  Whilst radio-loud quasars are only a small subset of the quasar population, the construction of radio-loud quasar samples is less prone to selection effects than are optical samples, since radio emission is unaffected by either intrinsic or extrinsic absorption due to dust. Fall \\& Pei (1993) discuss how dust within intervening galaxies may affect the observed evolution in optically selected samples of quasars.  Our survey technique involves optical identification of a large number of radio sources, followed by selection of the optically red, unresolved objects. The technique is similar to that used by Hook et al. (1995,1996) to find $\\sim 30$ $z>3$ quasars at higher radio flux densities, and complements the FIRST Bright Quasar sample (FBQS, Gregg et al. 1996).  The radio data are described in section 2 of this paper and the optical identification of the radio sample and selection of the spectroscopic sample are described in section 3. The spectroscopic observations are described in section 4, and the results are presented in section 5.  Finally, tables of redshifts, optical and radio data are given for the objects for which we have obtained spectra, along with the spectra and finding charts for the $z>3$ quasars. ",
        "conclusions": ""
    },
    "9804/astro-ph9804043_arXiv.txt": {
        "abstract": "We will report the observations of TeV gamma ray flares from Markarian 501 using Telescope Array Prototype. The observation were carried out continuously from the end of March to the end of July in 1997. The energy spectrum, and the time variation of the gamma ray intensities are shown. The intensity has been changed by the order of magnitude in this period and the possible quasi periodic oscillation of 12.7days were discovered. ",
        "introduction": "Markarian 501 (Mkn501) is an extragalactic BL-Lac type object at z=0.034. It is observed in radio, optical and X-ray bands, and can be characterized in flat spectrum radio, highly optically polarized and optically violently variable. BL-Lac type objects are considered to be a kind of AGN, which have a jet oriented to our line of the sight. Recently, Mkn501 (\\cite{Whipple-501,HEGRA-501}), Mkn421 (\\cite{Whipple-421,HEGRA-421,TA-421}), and 1ES 2344+514 (\\cite{Whipple-1ES}) have been identified as TeV gamma ray sources, and they are all BL-Lac objects. Gamma rays from Mkn501 around 100MeV to 10GeV was not observed by the CGRO EGRET detector(\\cite{EGRET}), however, emission in the TeV energy region were discovered by the Whipple telescope (\\cite{Whipple-501}) and confirmed by the HEGRA telescopes (\\cite{HEGRA-501}). Detailed study of gamma rays from AGN will give us information about the environment surrounding the huge black hole located at the center of the AGN, and the high energy phenomena and particle acceleration in the jet. The time scale of the intensity variations of the TeV gamma rays may explain the particle acceleration site is close to the black holes. TeV gamma rays from such extragalactic object interacts with the infrared background photons, and a cutoff of the energy spectrum is expected around 7-15TeV(\\cite{stecker}). There is an ambiguity in the prediction of the cutoff energy due to the uncertainty of the number density of infrared photons. In the discovery stage, the intensity in the TeV energy range was very small, corresponding to 8\\% of the Crab nebula flux, however, from Mar 1997 it increased dramatically and variated from 0.3 to 4 Crab flux. It is found that the variation of the intensities are larger in the TeV range than X-ray and other ranges(\\cite{Whipple-multi}). These flares are observed by Whipple , HEGRA, CAT, TACTIC and the Telescope Array Prototype (TAP) (\\cite{Whipple-501,HEGRA-501,CAT-501,TACTIC-501,TA-501}). The flares were continuously observed until the end of July 1997. Here we will report the details of the Mkn501 observations by TAP. ",
        "conclusions": "We obtained the differential energy spectrum of the gamma rays from Mkn501. It shows a possible cutoff feature above 5-7 TeV. This cutoff is predicted to be caused by the interaction with the infrared photons or the the limit of the electron acceleration in the jet. The cutoff energy of the gamma rays due to the infrared photon interaction from the distance of the Mkn501, z=0.034 is estimated 7 TeV to 15 TeV by \\cite{stecker}.  In the time variation of the gamma ray intensities, we found a periodicity, which may be a quasi periodic oscillation. The observed periodicity corresponds to 12.7 days or 25.5 days. The relationship of this periodicity with high energy phenomena around the massive black holes could be influenced by factors including the precession of the jet, and the rotation of the black hole(s). R. Protheroe (\\cite{Ray}) suggested the interaction of the shock wave and the helical structure of the jet may cause this observed type of periodicity."
    },
    "9804/astro-ph9804275_arXiv.txt": {
        "abstract": "We present JHK photometry and spectroscopy of RX J0019.8+2156.  The spectrum appears to be dominated by the accretion disc to at least 2.4 $\\mu$m, over any other source of emission.  We find Paschen, Brackett and He II lines strongly in emission, but no He I.  There are satellite lines approximately 850km s$^{-1}$ either side of the strongest, unblended hydrogen lines. These satellite lines may be the spectral signature of jets from the accretion disc.  \\begin{keywords} accretion, accretion disks -- ISM: jets and outflows -- stars: emission-line -- X-rays: stars -- infrared: stars -- stars: individual: RX J0019.8+2156  \\end{keywords} ",
        "introduction": "\\label{sec:intro}  The `supersoft' X-ray sources form a new and exciting but inhomogeneous class of X-ray objects. Though a few members of the class were known before ROSAT, it is this observatory which has characterised the class as objects with temperatures of 15 -- 80 eV and luminosities approaching the Eddington limit for a solar mass object, i.e.  $10^{36}$ -- $10^{38}$ erg s$^{-1}$. The model which best fits the observations has a white dwarf steadily or cyclically burning hydrogen-rich material accreted onto its surface at $\\sim 10^{-7}$ M$_{\\sun}$ yr$^{-1}$ (van den Heuvel et al. \\cite{vdheuvel92}).  van den Heuvel et al. (\\cite{vdheuvel92}) theorise that this material may be supplied by a near main-sequence A/F companion star, more massive than the white dwarf and in a close binary orbit, which is overflowing its Roche lobe.  Such systems are known as close binary supersoft sources (CBSS), for further discussion of the nature of the class see Kahabka \\& van den Heuvel (\\cite{kahabka97}).  35 supersoft X-ray sources have so far been identified within our galaxy, M31, the Local Group galaxy NGC 55, the LMC \\& SMC (listed in Kahabka \\& van den Heuvel \\cite{kahabka97}). Of these systems \\object{RX J0019.8+2156} (hereafter RX J0019) is the only northern hemisphere CBSS identified, so far, in our galaxy.  RX J0019 is a close binary with a 15$\\fh$85 period and both short-term and long-term optical variability, with orbital variability of $\\Delta V \\sim 0.5$ (Beuermann et al. \\cite{beuermann95}, hereafter B95; Greiner \\& Wenzel \\cite{greiner95}). UBVRI photometry at photometric maximum and minimum reveal a very blue continuum, in common with other CBSS.  The optical and UV spectra show Balmer $\\alpha$ and He II lines, especially $\\lambda\\lambda$1640,4686\\AA, in emission with enhanced blue wings, but no He I lines (B95).  Balmer $\\beta$ and $\\gamma$ lines exhibited P-Cygni type profiles. ",
        "conclusions": "\\begin{table} \\caption{Parameters of fits to HI lines using 3 Gaussians per line, where velocities are relative to the laboratory wavelengths.} \\begin{flushleft} \\begin{tabular}{cccc} \\hline Spectral&$V_{peak}$&FWHM&Peak flux\\\\ feature&(km s$^{-1}$)&(km s$^{-1}$)&(mJy)\\\\ \\hline P$\\gamma$-- & --790 $\\pm$ 197 & 542  $\\pm$ 441 & 0.8 $\\pm$ 0.4\\\\ P$\\gamma$  & --64  $\\pm$ 52  & 577  $\\pm$ 159 & 3.4 $\\pm$ 0.5\\\\ P$\\gamma$+ & +850 $\\pm$ 260 & 1026 $\\pm$ 614 & 1.0 $\\pm$ 0.3\\\\ \\hline P$\\beta$--  & --891 $\\pm$ 45  & 488  $\\pm$ 114 & 1.4 $\\pm$ 0.3\\\\ P$\\beta$   & --33  $\\pm$ 14  & 525  $\\pm$ 37  & 5.5 $\\pm$ 0.3\\\\ P$\\beta$+  & +838 $\\pm$ 44  & 606  $\\pm$ 114 & 1.7 $\\pm$ 0.2\\\\ \\hline Br$\\gamma$--& --931 $\\pm$ 163 & 386  $\\pm$ 430 & 0.3 $\\pm$ 0.2\\\\ Br$\\gamma$ & +2   $\\pm$ 33  & 531  $\\pm$ 85  & 1.8 $\\pm$ 0.2\\\\ Br$\\gamma$+& +832 $\\pm$ 167 & 556  $\\pm$ 438 & 0.4 $\\pm$ 0.2\\\\ \\hline \\end{tabular} \\end{flushleft} \\label{tab:jets} \\end{table}  No photospheric absorption lines are seen in the infrared spectrum taken at minimum light, indicating that either the underlying spectral type of the donor is not later than F, we are seeing the system at low inclination and irradiation is supressing absorption, or the accretion disc is dominating the emission.  Through the optical and IR to at least 2.4 $\\mu$m the orbital modulation of flux appears to be nearly constant, implying that they are dominated by the same component. The continuum is flatter than that expected for a stellar photosphere, suggesting the accretion disc is dominant.  Given the variability of the source over a variety of timescales (Greiner \\& Wenzel \\cite{greiner95}), simultaneous multiwavelength observations of RX J0019 would be necessary to properly determine the relative contribution of light sources, and consequently place constraints on the spectral class of the donor, inclination of and distance to the system.  One explanation of the satellite lines to P$\\gamma$, P$\\beta$ and Br$\\gamma$ is that they are the infrared signature of jets, similar to the Doppler shifted components of the Balmer $\\alpha$, $\\beta$ and He II $\\lambda$4686\\AA\\ lines of another supersoft source, \\object{RX J0513.9-6951} (hereafter RX J0513, Crampton et al. \\cite{crampton96}; Southwell et al.  \\cite{southwell96}), in which $v_{\\rm jet} \\cos i \\sim 3800$ km s$^{-1}$.  With this interpretation in mind we fitted each of these three composite lines profiles with three Gaussians. The details of these Gaussians are given in table \\ref{tab:jets}. The mean separation of the satellite lines from the core is $844 \\pm 76$ km s$^{-1}$, and their mean FWHM is $600 \\pm 220$ km s$^{-1}$.  Assuming the jet interpretation is correct, and the jets are perpendicular to an accretion disc which lies in the plane of the binary, the true velocity of the outflow is simply a function of the inclination angle ($v_{\\rm jet} = (844 / \\cos{i})$ km s$^{-1}$.). Following Shahbaz et al. (\\cite{shahbaz97}) we can also estimate the opening angle of the jet, $\\phi = \\sin^{-1} (\\Delta v/ v_{\\rm obs}\\tan{i})$, where $\\Delta v$ is the FWHM of the satellite lines and $v_{\\rm obs}$ the observed velocity.  Meyer-Hofmeister et al. (\\cite{meyer97}) model orbital modulation of the system with a best-fit inclination of 56$^{\\circ}$, corresponding to a outflow velocity of $\\sim 1500$ km s$^{-1}$ and an opening angle of 28$^{\\circ}$. Hachisu \\& Kato (\\cite{hachisu98}) model outbursts of the source with an inclination of 45$^{\\circ}$, corresponding to a velocity of $\\sim 1200$ km s$^{-1}$ and a very large opening angle of 45$^{\\circ}$.  An inclination angle of around 75$^{\\circ}$ would be required to derive the same velocity as the jet in RX J0513, and would imply a much more collimated jet ($\\phi \\sim 10^{\\circ}$). Alternatively the lower jet velocity could be intrinsic and reflect the lower escape velocity required from a lower mass white dwarf. Indeed Hachisu \\& Kato (\\cite{hachisu98}) argue for a mass of around 0.6 M$_{\\odot}$ for the white dwarf in RX J0019, around half that estimated for RX J0513 (Southwell et al. \\cite{southwell96}).  {\\em Note added in manuscript --} Just before submission we became aware of a related work by Becker et al.  (\\cite{becker98}). They have found transient satellite lines around hydrogen and He II lines in optical spectra of RX J0019. They also interpret the satellite lines to be the signature of jets and find $v_{jet}\\cos{i} \\sim$ 800 km s$^{-1}$, consistent with this work."
    },
    "9804/astro-ph9804039_arXiv.txt": {
        "abstract": "We present circumstantial evidence that the central region of the edge-on S0 galaxy NGC~4570, which harbors a 150~pc scale nuclear disc in addition to its main outer disc, has been shaped under the influence of a small ($\\sim 500$~pc) bar.  This is based on the finding of two edge-on rings, whose locations are consistent with the Inner Lindblad and Ultra Harmonic Resonances of a rapidly tumbling triaxial potential. Observed features in the photometry and rotation curve nicely correspond with the positions of the main resonances, strengthening the case for a tumbling bar potential. The relative blue colour of the ILR ring, and the complete absence of any detected ISM, indicates that the nuclear ring is made of relatively young ($\\lta 2$ Gyr) stars. We discuss a possible secular-evolution scenario for this complex multi-component galaxy, which may also apply to many other S0-galaxies with observed rings and/or multiple disc-components. ",
        "introduction": "High resolution images obtained with the Hubble Space Telescope (HST) have revealed the presence of bright stellar discs with very small scale lengths of $\\sim 20$~pc in a number of S0 galaxies (van den Bosch \\etal 1994; Lauer \\etal 1995; van den Bosch, Jaffe \\& van der Marel 1998).  These nuclear discs are photometrically and kinematically decoupled from the main (outer) discs which have an inner cut-off radius; the so-called Freeman type II discs (Freeman 1970). Such a multi-component structure suggests a complex formation scenario. The main question is whether this multi-component structure was imposed during the formation of the galaxy or whether it resulted from secular evolution.  Baggett, Baggett \\& Anderson (1996) argued that the inner-truncated discs could be linked with the presence of a bar. Since the nuclear regions of these S0s exhibit high line strength values (Fisher, Franx \\& Illingworth 1995), it has been suggested that they formed from radial inflow of pre-enriched gas. Small scale (nuclear) bars have been invoked as an efficient way to transport gas inwards and accumulate mass in the central regions of galaxies (Friedli \\& Martinet 1993; Wada \\& Habe 1995).  Bar-induced secular evolution could thus be a possible mechanism by which to form the observed nuclear discs. In this scenario the inner cut-off of the outer main disc would occur at the co-rotation radius of the bar, which drives gas into the central region where the nuclear disc forms. Unfortunately, there is a conspiracy in that nuclear discs can only be easily detected in nearly edge-on galaxies, whereas the photometric bar signatures (e.g., isophote twists) are confined to the more face-on systems.  In this paper we present the case of NGC~4570, an edge-on S0 galaxy with a nuclear stellar disc. Based on the study of the photometric and kinematical substructures in this galaxy (Section~2), we present evidence for the presence of a single pattern speed associated with a small bar (Section~3). This evidence is based mainly on the locations of two edge-on stellar rings. In Section~4 we discuss a secular evolution scenario for NGC~4570 under the influence of this bar potential. Conclusions are drawn in Section~5. ",
        "conclusions": "We have presented photometric evidence for the presence of two rings in the edge-on S0 galaxy NGC~4570. The main disc of this galaxy has an inner cut-off, inside of which a bright nuclear disc resides.  The locations of the rings are consistent with the Inner Lindblad and Ultra Harmonic resonances induced by a rapidly tumbling, triaxial potential with a pattern speed of $349 \\kmskpc$. The locations of the resonances correspond to local features in the morphology, colours and kinematics of NGC~4570. In particular, the outer disc has its inner cut-off close to CR, consistent with the picture that this cut-off originates from the bar (early-type bars generally end near CR, Combes \\& Elmegreen 1993).  The line strengths and colours of the central nucleus (inside the ILR) suggests a metal-rich and rather young stellar population. The ring at the ILR is bluer than its immediate surrounding, which together with the complete absence of any observed ISM, indicates that it consists of a relatively young ($\\lta 2$ Gyr) stellar population.  We have thus proposed a tentative scenario for the formation and secular evolution of NGC~4570. We have argued that a rapidly tumbling bar formed in NGC~4570 driving gas inwards thus forming part of the central structure (e.g., cusp and nuclear disc). Due to this central mass accumulation a strong ILR is created, and the bar then weakened. This was followed by the formation of an ILR gas ring which underwent subsequent star formation. The observed feature at the UHR also seems to be a genuine ring feature, since both observations and N-body simulations have shown that rings often form at the UHR as well. However, deeper photometric observations and a comparison of observed kinematics at the UHR with N-body simulations are required to confirm this hypothesis.  It is intriguing that no ISM is observed in NGC~4570, whereas the formation of rings clearly requires the presence of gas. However, estimates of the mass of the rings, do not require a large gas supply, but are consistent with them having formed from stellar mass loss of the older population of stars.  Several {\\it spiral} galaxies with double disc structure and/or multiple rings whose locations are consistent with the resonant radii of a single pattern speed are known: Emsellem \\etal (1996) presented the case of the Sombrero galaxy, M~104, an Sa which exhibits UHR and OLR rings as well as two distinct discs.  The Sb galaxy NGC 7217 harbors three rings (at ILR, UHR, and OLR), but lacks any photometric evidence for a bar (Athanassoula 1996).  Freeman type II discs and stellar rings have also been observed in a significant number of {\\it lenticular} galaxies: Seifert \\& Scorza (1996) detected rings or double disc structures in $60$\\% of their sample of 15 S0s. Although this sample has to be completed and higher spatial resolution data obtained (e.g., to search for nuclear discs and rings), this number is consistent with the percentage of more face-on disc galaxies which show photometric evidence for a bar or oval distortion (Combes 1995 and references therein). All this strongly suggests that double disc structures are common and most likely the result of bar driven secular evolution.  In this paper we have presented the case of NGC~4570, the first S0 galaxy known that shows the presence of {\\it two} stellar rings in addition to a double disc structure. We have discussed a secular evolution scenario to form this complex morphology, and which may also apply to other S0 galaxies with double disc structure and/or ring(s)."
    },
    "9804/nucl-th9804027_arXiv.txt": {
        "abstract": "Properties of dense nucleon matter and the structure of neutron stars are studied using variational chain summation methods and the new Argonne $v_{18}$ two-nucleon interaction, which provides an excellent fit to all of the nucleon-nucleon scattering data in the Nijmegen data base. The neutron star gravitational mass limit obtained with this interaction is 1.67\\msol. Boost corrections to the two-nucleon interaction, which give the leading relativistic effect of order $(v/c)^2$, as well as three-nucleon interactions, are also included in the nuclear Hamiltonian.  Their successive addition increases the mass limit to 1.80 and 2.20 \\msol. Hamiltonians including a three-nucleon interaction predict a transition in neutron star matter to a phase with neutral pion condensation at a baryon number density of $\\sim 0.2$ fm$^{-3}$. Neutron stars predicted by these Hamiltonians have a layer with a thickness on the order of tens of meters, over which the density changes rapidly from that of the normal to the condensed phase.  The material in this thin layer is a mixture of the two phases.  We also investigate the possibility of dense nucleon matter having an admixture of quark matter, described using the bag model equation of state. Neutron stars of 1.4\\msol do not appear to have quark matter admixtures in their cores. However, the heaviest stars are predicted to have cores consisting of a quark and nucleon matter mixture.  These admixtures reduce the maximum mass of neutron stars from 2.20 to 2.02 (1.91)~\\msol for bag constant $B = 200$ (122) MeV/fm$^3$.  Stars with pure quark matter in their cores are found to be unstable.  We also consider the possibility that matter is maximally incompressible above an assumed density, and show that realistic models of nuclear forces limit the maximum mass of neutron stars to be below 2.5\\msol. The effects of the phase transitions on the composition of neutron star matter and its adiabatic index $\\Gamma$ are discussed. ",
        "introduction": "The significant influence of nuclear forces on neutron star structure is by now firmly established by a large body of theoretical and observational evidence \\cite{st}. In the absence of these forces, the maximum possible mass of neutron stars composed of non-interacting neutrons is $\\sim 0.7$ solar masses (\\msol) \\cite{OV39}. Since most observed neutron star masses are above 1.3\\msol \\cite{nstarm}, they must be supported against gravitational collapse by pressure originating from nuclear forces. In the present work, we study neutron star structure using one of the most realistic models of nuclear forces currently available. A brief outline of previous calculations leading to this work is presented below.  Shortly after the discovery of pulsars, calculations of the equation of state (EOS) of neutron star matter with realistic models of the two nucleon interaction (NNI), obtained by fitting the nucleon-nucleon (NN) scattering data then available, were carried out using the lowest order constrained variational method \\cite{Pan71,BJ74}.  The results demonstrated that nuclear forces increase the mass limit of stable neutron stars beyond 1.4\\msol.  By the late 1970s it had become clear that the NNI alone could not account for the properties of nuclear matter or few-body nuclei. Variational \\cite{PW79} and Brueckner calculations \\cite{DW85}, including higher order cluster contributions, established that nuclear matter with realistic NNI saturates at too high a density. In addition, these interactions were known to underbind $^3$H. Ignoring the latter problem, plausible density dependent terms were added to the Urbana $v_{14}$ (U14) model of NNI \\cite{LP81a,LP81b} to reproduce the observed equilibrium properties of nuclear matter.  The resulting density dependent (U14-DDI) model of nuclear forces predicted stable neutron stars having masses up to 1.8\\msol \\cite{FP81,LRP93}.  Since nucleons are made up of quarks and have internal degrees of freedom, we can expect interactions among three (and perhaps four or more) nucleons, in addition to the NNI. The Urbana three nucleon interaction (TNI) models contain only two terms, with strengths fixed by the saturation density of nuclear matter and the binding energy of $^3$H.  Wiringa, Fiks and Fabrocini (WFF) \\cite{WFF88} used the U14 and the subsequent Argonne $v_{14}$ (A14) \\cite{WSA84} models of NNI, together with the Urbana VII (UVII) model of TNI, to study neutron star structure and obtained mass limits of 2.19\\msol and 2.13\\msol with the U14+UVII and A14+UVII, respectively. They also found that pure neutron matter (PNM) undergoes a transition to a phase having spin-isospin order, attributed to neutral pion condensation, at a density of $\\sim 0.2$ fm$^{-3}$ with the A14+UVII, but not with the U14+UVII. Neither of these models results in a phase transition in symmetric nuclear matter (SNM), which is composed of equal numbers of neutrons and protons.  In the early 1990s the Nijmegen group \\cite{SKR93} examined carefully all NN scattering data at energies below 350 MeV published between 1955 and 1992.  They extracted 1787 proton-proton (pp) and 2514 proton-neutron (np) ``reliable'' data, and demonstrated that these data determine all NN phase shifts and mixing parameters quite accurately. The NNI models which fit this Nijmegen data base with a $\\chi^2/N_{data} \\sim 1$ are called ``modern''.  These include the Nijmegen models \\cite{SKT94}: Nijmegen I, II and Reid-93, the Argonne $v_{18}$ \\cite{WSS95}, denoted here by A18, and the CD-Bonn \\cite{MSS96}.  In order to fit the pp and np data simultaneously and accurately, these models include a detailed description of the electromagnetic interactions and terms that break the isospin symmetry of nuclear forces. All include the long range one-pion exchange potential, but follow different treatments of the intermediate and short range parts of the NNI. The differences among the predictions of these models for the properties of many-body systems are much smaller than those among the predictions of older models, presumably because all modern potentials accurately fit the same scattering data. For example, the $^3$H binding energies predicted by the modern Nijmegen models and A18 are between -7.62 to -7.72 MeV \\cite{FPS93}, while that of CD-Bonn is -8.00 MeV \\cite{MSS96}.  The difference between these results and the experimental value of -8.48 MeV is used to fix one of the parameters of Urbana TNI models.  Detailed studies of the energies of dense nucleon matter were carried out recently by Engvik {\\em et al.} \\cite{Eng97} using all the modern models of NNI and the lowest order Brueckner (LOB) method. According to these studies, the results with the modern potentials are all quite similar up to densities relevant to neutron stars.  For example, the energies predicted with the LOB method for the energy of neutron matter at 5$\\rho_0$, where $\\rho_0 = 0.16$ fm$^{-3}$ is the equilibrium (saturation) density of nuclear matter, range from 80 to 93 MeV per nucleon.  The spread of 13 MeV in these energies is small compared to the possible errors in the LOB method and the expected contributions of TNI at this density.  This model independence results from the fact that the mean interparticle distance at $\\rho \\sim 5\\rho_0$ is greater than 1~fm, and the predicted matter energy is therefore not sensitive to the details of the interaction at $r_{ij} < 1$~fm.  In this paper we study the structure of neutron stars with the A18 model using variational chain summation (VCS) methods, which  hopefully include all leading many-body correlation effects. The Urbana model IX (UIX) \\cite{Pud95} is used to estimate the effect of TNI.  Previous studies of nucleon matter with A18 and UIX interactions have indicated the possibility of a transition to a neutral pion condensed phase for both PNM and SNM \\cite{AP97}.  The effects of such a transition on the structure of neutron stars are studied here in detail.  The effect of relativistic boost corrections \\cite{FPF95} to the A18 interaction is also examined. At high densities, we consider the possibility of matter becoming maximally incompressible, as well as that of a transition to mixed phases of quark and nucleon matter \\cite{Gle92,HPS93}.  The relativistic mean field (RMF) approximation \\cite{Wal74} has been used in many studies of high density matter and neutron stars. There exists a vast amount of literature on this topic, some of which has been reviewed by Glendenning \\cite{Gle97}.  While the RMF approximation is very elegant and pedagogically useful, it is not valid in the context of what is known about nuclear forces, which is the theme of this work. For example, using the meson parameters of the CD-Bonn S-wave potentials in the RMF approximation leads to unbound SNM. (In order to accurately fit the NN scattering data, the phenomenological scalar meson parameters in the CD-Bonn model are allowed to depend on partial wave quantum numbers.) At a density of $\\rho_0$, for SNM, the RMF approximation yields an energy per nucleon of $\\sim +20$ MeV, while the LOB method gives $\\sim -18$ MeV.  In the mean field Hartree approximation, which is implicit in the RMF calculations, the A18 NNI gives energies per nucleon of $\\sim$ +30 (+37)~MeV at $\\rho_0$ and +155 (+204)~MeV at 5$\\rho_0$ for SNM (PNM), while the variational calculations presented here give $\\sim$ -18 (+12) and +25 (+88) MeV, respectively.  The main problem is that the mean field approximation for meson fields is only valid for $\\mu \\bar{r}\\ll 1$, where $\\mu$ is the inverse Compton wave length of the meson and $\\bar{r}$ is the mean interparticle spacing. Over the 1-5 $\\rho_0$ density range, $\\bar{r}$ estimated using a body centered cubic lattice ranges from 2 to 1.2 fm.  Thus $\\mu \\bar{r}$ is in the range 1.4 to 0.8 for the pion and 7.8 to 4.7 for vector mesons. The mean field approximation is not applicable, since these values are obviously far from being much smaller than one.  The RMF approximation can be based on effective values of the coupling constants that take into account the correlation effects. However, these coupling constants then have a density dependence, and a microscopic theory is needed to calculate them.  This paper is organized as follows.  Section II contains a summary of the non-relativistic calculations with A18 and A18+UIX models of nuclear forces, while section III describes the calculations including the relativistic boost interaction, denoted by $\\delta v$, without and with the TNI model UIX$^*$.  The beta equilibrium of neutron star matter is discussed in section IV and results for neutron star structure are presented in section V, where we also discuss the effects of the possible transition to mixed nucleon and quark matter phases.  The adiabatic index and sound velocities in neutron star matter are given in section VI, and conclusions are presented in section VII. ",
        "conclusions": "The results obtained with the modern A18 NNI for the gross properties of neutron stars, such as the mass limit and the radius of 1.4\\msol~stars, are not very different from those obtained from the earlier U14 and A14 models \\cite{WFF88}.  The main difference is that with A18 and A14 NNI there is the indication of the neutral pion condensation phase transition, while there is none with U14.  However, the fact that A18 provides an exact fit to the NN scattering data while U14 does not weighs against the U14 results.  As is well known, the cooling of neutron stars is accelerated by pion condensation \\cite{Pet92}. The detailed mechanism associated with the spin-ordered condensate that we find is as yet unexplored. In addition, we estimate that due to this phase transition, there may be regions in the star where the matter density changes rapidly over a ten meter distance scale.  The matter in this thin layer of the star can have interesting structures on a ten fermi length scale.  We find that inclusion of the relativistic boost correction, $\\delta v$, to the NNI increases the mass limit from 1.67\\msol to 1.8\\msol without any TNI, while it is reduced from 2.38\\msol to 2.20\\msol upon inclusion of the Urbana models of TNI.  The reduction occurs because the TNI needed to fit nuclear data have weaker repulsive parts after including $\\delta v$.  For the same reason, the effect of Urbana models of TNI on the mass limit is diminished from 0.71\\msol to 0.4\\msol when the $\\delta v$ is included.  Note that these models of TNI have only two terms; the long range two-pion exchange term, and a short range term with no spin-isospin dependence.  Their strengths are determined from the density of SNM and the triton energy. However the triton, which has isospin 1/2, is insensitive to the interaction between three neutrons.  Improved models of TNI must consider data such as the binding energy of $^8$He, which are sensitive to the interaction between three neutrons.  If the effective value of the bag constant B is larger than 122 MeV/fm$^3$ it appears that only the heaviest neutron stars may have small drops of quark matter in their interior.  The quark composition of these drops is similar to that of aggregates of $\\Sigma^-$ hyperons.  If the interaction of $\\Sigma^-$ with dense nucleon matter is not repulsive, there may exist $\\Sigma^-$ hyperons in nucleon matter at densities below the threshold for the appearence of quark drops.  It is necessary to build $N\\Sigma^-$ and $NN\\Sigma^-$ interaction models to more fully explore this possibility.  Using maximally incompressible matter at high densities we find that the upper limit for the maximum mass consistent with nuclear data is $\\sim$2.5\\msol ; this is not far from the prediction of 2.21\\msol of our A18+$\\delta v$+UIX$^*$ model.  The lower limit for the maximum mass is more difficult to establish due to the unknown interactions between hyperons and nucleons.  It may be as low as 1.74 $M_{\\odot}$ if the quark bag constant B has a value of $\\sim 122$ MeV/fm$^3$, and if the three neutron interaction is not as repulsive as the Urbana model IX at high densities, even without admixture of hyperons in matter.  Recently several authors \\cite{nx1,nx2,nx3} have argued that there are indications of the existence of neutron stars with $M \\sim 2M_{\\odot}$.  If these are confirmed, then models without TNI will be ruled out. However, such a possibility is still the subject of active debate \\cite{nx4}."
    },
    "9804/astro-ph9804259_arXiv.txt": {
        "abstract": "We analyse the stochastic background of gravitational radiation emitted by a cosmological population of core-collapse supernovae. The supernova rate as a function of redshift is deduced from an observation-based determination of the star formation rate density evolution. We then restrict our analysis to the range of progenitor masses leading to black hole collapse. In this case, the main features of the gravitational-wave emission spectra have been shown to be, to some extent, independent of the initial conditions and of the equation of state of the collapsing star, and to depend only on the black hole mass and angular momentum. We calculate the overall signal produced by the ensemble of black-hole collapses throughout the Universe, assuming a flat cosmology with vanishing cosmological constant. Within a wide range of parameter values, we find that the spectral strain amplitude has a maximum at a few hundred Hz with an amplitude between $10^{-28}$ and $10^{-27}$ Hz$^{-1/2}$; the corresponding closure density, $\\Omega_{GW}$, has a maximum amplitude ranging between $10^{-11}$ and $10^{-10}$ in the frequency interval $\\sim 1.5-2.5$ kHz. Contrary to previous claims, our observation-based determination leads to a duty cycle of order $0.01$, making our stochastic backgound a non-continuous one. Although the amplitude of our background is comparable to the sensitivity that can be reached by a pair of advanced LIGO detectors, the characteristic shot-noise structure of the predicted signal might be in principle exploited to design specific detection strategies. ",
        "introduction": "According to the theory of general relativity, gravitational waves are expected to be produced in such a large variety of astrophysical and cosmological phenomena that it is plausible to guess that our universe is pervaded by  a background of gravitational radiation. Depending on their origin, different contributions to this background will exhibit peculiar spectral properties, the features of which it is interesting to investigate in view of a possible detection by future gravitational antennas.  Gravitational waves of cosmological origin could be the result of processes that developed in the very early Universe. Due to the extremely small graviton cross-section, the spectral properties of such a relic radiation should have been retained with no substantial alteration until today, and, if observed, would provide information on the physical conditions of the epochs when these waves were produced (see, for a recent review, Maggiore 1998).  Further contributions to the stochastic background of gravitational waves are  of astrophysical nature and their generation  dates back to more recent epochs, when galaxies and stars started to form and evolve. For instance, the background generated by the radiation emitted  from rotating neutron stars in the Galaxy has been considered in (Giampieri 1997; Postnov 1997) and it has been shown that, by taking into account the uncertainty with which the ellipticity of neutron stars is known, $h^2\\Omega_{GW}(\\nu=100~\\mbox{Hz})$ is in the range $10^{-15} - 10^{-9}$, where \\op \\Omega_{GW}\\cl is the closure energy density of gravitational waves per logarithmic frequency interval, \\be \\Omega_{GW}=\\frac{1}{\\rho_c}\\frac{d\\rho_{GW}}{d\\log{\\nu}}, \\ee with \\op \\rho_{c}= 3H_0^2\\,/\\,8\\pi G\\cl the critical density.\\footnote{ Throughout this paper we shall write the Hubble constant as\\[H_0=h \\times 100\\, \\mbox{km}~\\mbox{s}^{-1}\\mbox{Mpc}^{-1}.\\]} A strategy for detecting these signals with one interferometric antenna, by squaring the detector amplitude and searching for a sidereal modulation has been proposed in (Giazotto, Bonazzola \\& Gourgoulhon 1997; Giampieri 1997). Further studies have estimated the intensity of the stochastic background produced by the galactic merging of unresolved binary white dwarfs (Postnov 1997), and subsequently extended to the extragalactic white dwarf's merging for different  cosmological models (Kosenko \\& Postnov 1998). The considered range  of frequency is \\op 10^{-3} - 10^{-2} \\, \\mbox{Hz}.\\cl   In this paper,  we study the background of gravitational waves emitted by sufficiently massive stars that, having reached the final stages of their evolution, collapse to form a black hole. We consider only collapses to a black hole because numerical studies have shown that the spectrum of the gravitational energy emitted in these processes exhibits some distinctive features that are to some extent independent of the initial conditions and of the equation of state of the collapsing star, and that depend only on the black hole mass and on its angular momentum (see, for a recent review, Ferrari \\& Palomba 1998). Conversely,  the simulations of gravitational collapse to a neutron star available  in the literature indicate that the energy spectrum which is produced strongly depends on the equation of state of the collapsing star, and no distinctive features can easily be extracted to model the process.  The  energy spectrum which we will use to model the collapse to a black hole is that obtained by Stark \\& Piran (1985, 1986), who studied the evolution of the axisymmetric collapse of polytropic, rotating stars, by a  fully relativistic numerical simulation.  In order to evaluate the contribution that each burst of gravitational radiation emitted in a collapse occurred in recent or past epochs gives to the stochastic background, the knowledge of the  star formation rate as a function of the cosmological redshift is essential. This function has been calculated in the framework of the hierarchical theory of galaxy formation (Cole et al. 1994; Kauffmann, White \\& Guiderdoni 1993; Navarro, Frenk \\& White 1996). However, in this paper we prefer to base our calculation on the star formation rate which can be deduced from observations, being  aware of the fact that, since these observations are reliable only up to redshifts of order \\op z\\sim 4-5,\\cl we may overlook the possibility that a big burst of star formation may have occurred in a far past (Rosi \\& Zimmermann 1976, Bertotti \\& Carr 1980, Bond \\& Carr 1984). Following  a recently proposed approach (Madau, Pozzetti \\& Dickinson 1997), the star formation rate history that we have considered is based on a very simple stellar evolution model which assumes a time dependent star formation rate and a constant initial mass function. The model is directly inferred from the impressive collection of recent observations of the rest frame UV-optical emission from star forming galaxies at redshifts ranging from zero up to $z \\simeq 4-5$ (Steidel et al. 1996; Madau et al. 1996; Lilly et al. 1996; Connolly et al. 1995; Ellis 1997), and assumes a Salpeter initial mass function, a flat cosmology with zero cosmological constant and $h=0.5$. The resulting star formation history has been extensively investigated and proves to correctly reproduce, within the experimental errors, the evolution of the galaxy emission properties observed in other wavebands. Furthermore, it  is consistent with the indications of QSO absorption lines and metallic clouds observations (Ellis 1997; Madau, Pozzetti \\& Dickinson 1997; Lanzetta, Yahil \\& Fernandez-Soto 1996; Pei \\& Fall 1995).\\\\  The plan of the paper is as follows. In Section 2 we briefly review the current state of observations of the global star formation in field galaxies and present the model for the star formation rate evolution that we have adopted in our analysis; in Section 3 we compute the  rate of collapses and the duty cycle of the process. In Section 4 we describe the  energy spectrum computed by Stark and Piran, which we use to model the gravitational emission of  single events. In Section 5 we derive the expression for the spectral energy density and for \\op \\Omega_{GW},\\cl and discuss the spectral properties of this background. In Section 6 we analyze the dependence of our results on the assumed model parameters. Finally, in Section 7 we present a preliminary statistical analysis of the signal. A discussion of  the assumptions on which our results are based and on  further developments of the present work will be given in the concluding remarks. ",
        "conclusions": "From the analysis carried out in this paper it emerges that a reliable estimate of the stochastic background of gravitational waves produced by a cosmological population of astrophysical sources cannot set the important effect of SFR evolution aside.  The extraordinary advances attained in observational cosmology are leading to a coherent picture for the global star formation history in field galaxies that can be used to infer the redshift evolution of the rate of gravitational wave sources.  More importantly, the resulting gravitational signal appears to be insensitive to the uncertainties that still trouble the high redshift SFR observations (as a consequence of the uncertain amount of dust obscuration), as the main contribution to the resulting $\\Omega_{GW}$ and $\\sqrt{S_h}$ comes from low-to-intermediate redshift sources.  We have restricted our investigation to a specific cosmological background model with $\\Omega_0=1, \\Lambda=0, H_0=50\\,\\mbox{km}\\,\\mbox{s}^{-1} \\mbox{Mpc}^{-3}$. We plan to extend the analysis to different choices of background cosmology in a subsequent paper. We can anticipate that, if we assume that the cosmological constant is zero, the SFR density decreases with $\\Omega_0$. At the same time, the geometrical effect on the comoving volume element increases the rate of black hole collapses compared to a flat background, although the decrease of the single event spectra with the square of the inverse luminosity distance partly cancels the effect. Altogether we would then expect a modest increase of amplitude of the resulting gravity-wave background.  The spectral properties of the background of gravitational waves discussed in this paper depend on the energy spectrum  we have chosen as representative of the process of gravitational collapse to a black hole. As mentioned in Section 4, it was derived by Stark \\& Piran (1985, 1986) and it refers to a simplified model of star, namely a naked compact core with a polytropic equation of state. At present, detailed information on the gravitational radiation emitted in collapses of stars with more realistic equations of state, and derived by a fully relativistic approach as in the Stark-Piran model, are not available.  As far as the ``naked core\" assumption is concerned, it should be mentioned that detailed studies on the structure of stars that evolve into Type II supernovae, show that quite a general pre-collapse configuration is composed by a degenerate iron  core surrounded by layers of matter, remnants of the earlier nuclear burning processes  (Woosley \\& Weaver 1995) However, these layers  are not expected to contribute significantly to the emitted radiation, because they do not participate in the core bounce (see for example the discussion in Section VII of Seidel 1991).  The preminent features of the Stark-Piran energy spectrum are (see Figure 3): \\begin{itemize} \\item The presence of a peak at the frequency of the quasi-normal modes of the formed black hole. \\item  The dependence of the efficiency on the fourth power of the angular momentum of the collapsing star. \\item  The presence of a further pronounced peak at a lower frequency, if the angular momentum is sufficiently high. \\end{itemize} It is interesting to ask to what extent are these features general, and how do they reflect in the structure of the quantities we compute, for instance  the strain amplitude, which is of direct interest for gravitational detectors. Gravitational signals emitted during the collapse of naked cores to a black hole have been studied in several papers with a perturbative approach:  a spherically symmetric collapsing star is perturbed axisymmetrically, and the emitted radiation is computed by solving the linearized equations (see Ferrari \\& Palomba 1998, for a recent review). The amount of  gravitational energy estimated by the perturbative approach is obviously lower than  that predicted by Stark and Piran. However, the presence of the peak of the quasi-normal modes can be considered as a general feature, unless the dynamics of the fluid is dominated by very strong internal pressure or by very long bounces, which basically slow down the collapse to such an extent that the formed black hole has no residual mechanical energy to radiate away in its quasi-normal modes (Seidel 1991). The ``two-peak\" structure of the energy spectrum can be mantained also if the star is non-rotating (Seidel 1991), if bounces occur during the collapse. The dependence of the emitted energy on the fourth power of the angular momentum is found also in the perturbative approach (Cunningham, Price \\& Moncrief 1980). In all cases the frequency scales as the inverse of the black hole mass, and the energy radiated as the second power of the mass. Thus, the features of the energy spectrum  we use to model each single event are likely to reasonably represent a generic situation.  The strain amplitude we computed from the Stark-Piran energy spectrum is shown in Figure 5. It exhibits a peak at some frequency which is the reminiscent of the peak of the quasi-normal modes of the formed black holes. Since we do not know  the angular momentum distribution of the black holes, each  curve in that figure is obtained by assuming that all black holes that form have the same value of the rotation parameter, and the location of the peak of \\op\\sqrt{S_h} \\cl depends on that value. We can summarize the above discussion by saying that the strain amplitude associated to the background under consideration should present a peak at a frequency between $\\sim$ 230 and 340 Hz, whereas the expected maximum amplitude should be of order \\op 1-8 ~\\times 10^{-28}.\\cl However, since our results are based on the assumption that each collapse takes place axisymmetrically the maximum strain amplitude may be underestimated: non-axisymmetric collapse would be more efficient in producing gravitational waves. The strain amplitude plotted in Figure 5 was derived for  the following set of parameters: $\\alpha=0.1$, $M_p=25\\msun$ and $M_u=125\\msun$. The dependence on different choices for the progenitor mass range ($M_p$, $M_u$)  was considered in section 6. The results are shown in Figure 8 and in Table 2. For a lower value of the upper mass cutoff, the frequency of the maximum is shifted by a factor $\\sim 2.4$ towards higher frequencies whereas the maximum amplitude is lowered by a factor $\\sim 2.2$. Similarly, Figure 9 shows the effects of variating the fraction of the progenitor mass which participates to the collapse ($\\alpha$). We can consider these curves as an indication of the uncertainties introduced by the lack of consistent results in the literature. However, reasonable values of the parameters $\\alpha$, $M_p$ and $M_u$ should not be too far from the ranges we have adopted. Thus, we may conclude that, according to our study, the spectral strain amplitude produced by the background of supernova explosions leading to black hole formation should have a maximum at a few hundred Hz with an amplitude ranging between \\op 10^{-28}\\cl and \\op 10^{-27}\\cl Hz$^{-1/2}$. The corresponding closure density, \\op \\Omega_{GW},\\cl has a maximum amplitude ranging between \\op 10^{-11}\\cl and \\op 10^{-10}\\cl in the interval \\op \\sim 1.5-2.5\\cl kHz. This is beyond the sensitivity that can be obtained by cross-correlating the outputs of two gravity-wave detectors of the first generation experiments (VIRGO-VIRGO, LIGO-LIGO). However, for a constant frequency spectrum the advanced LIGO detector pair aims at $\\Omega_{GW} \\sim 2 \\times 10^{-10}$ (Allen \\& Romano 1998), a level comparable to our signal.  Finally, we have made a preliminary statistical analysis of our background, e.g. in terms of the probability that the signal is above a given threshold. More details will be considered in a forthcoming paper. Let us stress, once again, that the smallness of the duty cycle, which makes our stochastic background a non-continuous one, implies a non-trivial statistics of the signal arrival times which will be a distinctive feature of the process that can be in principle exploited to design optimal detection strategies."
    },
    "9804/astro-ph9804290_arXiv.txt": {
        "abstract": "We examine the K shell emission lines produced by isothermal and simple multiphase models of the hot gas in elliptical galaxies and galaxy clusters to determine the most effective means for constraining the width of the differential emission measure ($\\xi(T)$) in these systems which we characterize by a dimensionless parameter, $\\sigma_\\xi$. Comparison of line ratios of two-temperature $(\\sigma_\\xi\\ll 1)$ and cooling flow $(\\sigma_\\xi\\sim 1)$ models is presented in detail. We find that a two-temperature model can approximate very accurately a cooling flow spectrum over 0.5-10 keV.  We have re-analyzed the {\\sl ASCA} spectra of three of the brightest galaxy clusters to assess the evidence for multiphase gas in their cores: M87 (Virgo), the Centaurus cluster, and the Perseus cluster. K$\\alpha$ emission line blends of Si, S, Ar, Ca, and Fe are detected in each system as is significant Fe K$\\beta$ emission.  The Fe K$\\beta$/K$\\alpha$ ratios are consistent with optically thin plasma models and do not suggest resonance scattering in these systems. Consideration of both the ratios of H-like to He-like K$\\alpha$ lines and the local continuum temperatures clearly rules out isothermal gas in each case. To obtain more detailed constraints we fitted plasma models over 1.6-9 keV where the emission is dominated by these K shell lines and by continuum. In each case the {\\sl ASCA} spectra cannot determine whether the gas emits at only two temperatures or over a continuous range of temperatures as expected for a cooling flow. The metal abundances are near solar for all of the multiphase models.  We discuss the implications of these results and examine the prospects for determining the temperature structure in these systems with upcoming X-ray missions. ",
        "introduction": "\\label{intro}  There is mounting evidence that the X-ray emission in elliptical galaxies and galaxy clusters originates from multiphase plasma.  The first important evidence for multiphase gas in ellipticals and clusters came from analysis of emission lines in M87 and the Perseus cluster.  Using ratios of several blends of Fe L shell lines obtained from the Focal Plane Crystal Spectrometer (FPCS) on the {\\it Einstein} Observatory, Canizares et al.  \\shortcite{crc82} found that the hot gas in M87 emits over a large range of temperatures: from $\\sim 2-4\\times 10^6$ K to $\\sim 3\\times 10^7$ K. The strength of the Fe XVII line measured from FPCS data of the Perseus cluster also suggests a multiphase plasma (Canizares, Markert, \\& Donahue 1988).  Recently, data from the {\\sl ROSAT} \\cite{rosat} and {\\sl ASCA} \\cite{tanaka} satellites have provided new evidence for multiphase hot gas. Broad-band spectral analysis of {\\sl ASCA} data of clusters (Fukazawa et al. 1994; Fabian et al. 1994a), the brightest ellipticals (Buote \\& Fabian 1998; Buote 1999a), and of the brightest poor galaxy groups \\cite{bgroup} indicates that the hot gas in the central regions of these systems consists of at least two temperature components. The equivalent widths and ratios of K$\\alpha$ line blends of Si and S alone strongly favor multitemperature models of the {\\sl ASCA} data of the brightest ellipticals \\cite{bmulti}.  Complementary evidence for multiphase gas in clusters is provided by the analysis of imaging data from {\\sl ROSAT} \\cite{rosat} of cooling flow clusters which infers mass deposition rates which vary with radius approximately as $\\dot{M}\\propto r$ (e.g. Peres et al. 1998). This type of distributed mass deposition is the signature of multiphase cooling flows (for a review see Fabian 1994). Multiphase cooling flows are also consistent with the positive temperature gradients observed in the centers of the brightest elliptical galaxies with {\\sl ROSAT} (e.g. Forman et al. 1993; David et al. 1994; Trinchieri et al. 1994; Rangarajan et al. 1995; Irwin \\& Sarazin 1996; Jones et al. 1997; Buote 1999a).  It is essential that the temperature structure be accounted for in order to deduce correctly the properties of the hot gas in these systems.  The metal abundances are especially sensitive to the assumed temperature structure.  For example, analyses of the X-ray spectra of elliptical galaxies where the hot gas is assumed to be isothermal always obtain very sub-solar metallicities, $Z\\la 0.3Z_{\\sun}$, whereas when multi-temperature models are considered the metallicity of ellipticals becomes consistent with solar (see Buote \\& Fabian 1998 and Buote 1999a). A similar situation prevails for poor groups of galaxies \\cite{bgroup}.  Although this ``Fe bias'' is unimportant for clusters of galaxies with temperatures in excess of $\\sim 3$ keV, if an isothermal gas is assumed for a cooling-flow cluster (or elliptical or group) then the Si/Fe and S/Fe ratios can be significantly over-estimated (see the appendix in Buote 1999b). Hence, in order to obtain reliable measurements of the heavy element abundances crucial to interpreting models of the formation, evolution, and enrichment of the hot gas in ellipticals, groups, and clusters (e.g. David, Forman, \\& Jones 1991; Loewenstein \\& Mathews 1991; Ciotti et al. 1991; Loewenstein \\& Mushotzky 1996; Renzini 1997; Brighenti \\& Mathews 1999), it is necessary to know the true temperature structure of the hot gas.  Knowledge of the temperature structure is also required to understand the contribution of discrete sources to the hard X-ray emission in elliptical galaxies.  For example, if the hot gas is assumed to be isothermal then the hard spectral component measured by {\\sl ASCA} for the elliptical galaxy NGC 3923 is inconsistent with the expectation that discrete sources follow the same profile as the optical light (Buote \\& Canizares 1998; see also Pellegrini 1999).  Finally, when multiphase cooling flows are accounted for in X-ray observations of clusters the scatter in the X-ray luminosity-temperature relationship is significantly reduced (e.g. Fabian et al. 1994b).  The Fe L shell emission lines from $\\sim 0.7$ - 1.4 keV generally provide the strongest constraints on the temperature structure of ellipticals, groups, and clusters. There has been considerable progress in improving the accuracy of the theoretically computed Fe L shell transitions (e.g Liedahl et al.  1995), but they are still among the most uncertain transitions implemented in the currently available plasma codes (e.g. Savin et al. 1999). There is reason for some confidence in the codes since recent empirical studies have concluded that the errors in the Fe L transitions are not overly serious (Hwang et al. 1997; Buote \\& Fabian 1998). (However, see the recent papers by Savin 1999 and Savin et al. 1999 for discussions of the inaccuracies in the dielectronic recombination rates used by the popular plasma codes which affect all types of emission lines.)  Nevertheless, it is important to use lines from the more-accurate K shell transitions to search for evidence of multiphase gas to complement the information from Fe L lines.  To obtain the most accurate constraints on the emission measure of a multiphase plasma it is necessary to obtain independent measurements of the temperature(s) over the entire energy range of significant emission. It is a happy coincidence that the K shell lines of many abundant elements appear throughout the X-ray wave band (0.1-10 keV).  Line ratios have important advantages over broad-band spectral fitting for probing the temperature structures of ellipticals and clusters. Firstly, in an isothermal plasma using ratios of lines of the same element avoids uncertainties in the relative abundances of the elements; there is a small dependence on metallicity in a multiphase cooling flow (section \\ref{cfs}).  Secondly, ratios of lines nearby in energy are essentially independent of Galactic or any intrinsic absorption. Thirdly, since the strength of a line can be computed using a local continuum measurement, line ratios are insensitive to any poorly constrained power-law or bremsstrahlung spectral components.  For a multiphase {\\it coronal} plasma \\footnote{The assumptions of the standard coronal model are discussed in, e.g., Sarazin \\shortcite{sarazin} and Mewe \\shortcite{mewe91}.}  the luminosity of an emission line arising from a transition from an upper state $u$ to a lower state $l$ of an element $X$ in ionization state $r$ may be expressed as, \\begin{equation} L_{ul}(X_r) = \\int \\Lambda_{ul}(X_r,T)\\xi(T)\\,dT, \\label{eqn.line} \\end{equation} where $\\Lambda_{ul}(X_r,T)$ is the plasma emissivity of the line at temperature $T$ which incorporates the abundance of element $X$, the relative fraction of ionization state $r$, and other quantities depending only on the atomic physics. (The integral extends over the emitting volume $V$.) The quantity, \\begin{equation} \\xi(T)\\equiv n^2_e(T)\\frac{dV}{dT}, \\label{eqn.dem} \\end{equation} is called the ``differential emission measure'' and $\\xi(T)dT$ the ``emission measure'' at temperature $T$. The emission measure weights $\\Lambda$ as a function of temperature in equation (\\ref{eqn.line}) and thus effectively characterizes the temperature structure of the plasma.  Notice that different phases $T$ in general can have different electron number densities $n_e$.  It has been known for some time \\cite{cb} that precise constraints on $\\xi(T)$ from inversion of equation (\\ref{eqn.line}) are difficult to obtain in practice.  (For a recent discussion of this topic see, e.g., Judge, Hubeny, \\& Brown 1997.)  That is, for a given line typically $\\Lambda_{ul}(X_r,T)$ is significant over temperature ranges of factors of $\\sim 2$-$5$, which means that the plasma emissivity acts as a smoothing function which obscures the detailed shape of $\\xi(T)$. Moreover, small uncertainties in the data are amplified by the inversion. Given these difficulties, it may only be possible to obtain relatively crude constraints on the shape of $\\xi(T)dT$ for ellipticals and clusters from inversion of equation (\\ref{eqn.line}); e.g., width, skewness, etc \\footnote{See, e.g., Kaastra et al. \\shortcite{kas} and Schmitt et al.  \\shortcite{schmitt} for recent constraints on $\\xi(T)$ of stellar coronae.}.  The primary theoretical goal of this paper is to identify the best X-ray emission line ratios for determining the gross features of $\\xi(T)$ of ellipticals and clusters with the specific objective of distinguishing between single-temperature, two-temperature, and cooling-flow spectra. (Previous discussions of important temperature sensitive lines for clusters in the context of a different set of models may be found in Bahcall \\& Sarazin \\shortcite{bs77}, Sarazin \\& Bahcall \\shortcite{sb77}, and Canizares et al. \\shortcite{crc82}. Wise \\& Sarazin \\shortcite{ws93} discuss the radial variation of the shapes of emission lines as diagnostics for cooling flow models.) To facilitate quantitative comparison of the multiphase structure of different ellipticals and clusters we introduce a quantity, $\\sigma_\\xi$, the ``multiphase strength''.  We focus on K shell emission lines since they are generally the strongest lines of ellipticals and clusters and the atomic physics is relatively simple and so the available plasma codes should be most accurate for these transitions.  Because the Fe L shell lines occurring at energies $\\sim 0.7-1.4$ keV are strong for ellipticals and clusters and are powerful as temperature diagnostics, we also examine them even though at the present time the accuracy of the plasma codes is less certain for these transitions (see above). The line ratios allow simple yet effective discrimination between popular spectral models of ellipticals and clusters and serve as an important guide to interpreting results of $\\xi(T)$ obtained from inversion of equation (\\ref{eqn.line}).  The paper is organized as follows. In Section \\ref{theory} we analyze theoretical line ratios in single-phase (Section \\ref{iso}) and multiphase (Section \\ref{multi}) models. The multiphase strength is introduced in Section \\ref{mps} and it is used to elucidate the temperature structure of simple multiphase models in Section \\ref{examples} and of cooling flow models in Section \\ref{cfs}. We discuss in detail how to distinguish a two-temperature plasma from a cooling flow in Section \\ref{2tcf}.  In Section \\ref{asca} we compute and analyze line ratios of {\\sl ASCA} data of M87, the Centaurus cluster, and the Perseus cluster. Simulated observations with {\\sl Chandra}, {\\sl XMM}, and {\\sl ASTRO-E} are considered in Section \\ref{sims}.  Finally, we comment on recent evidence for resonance scattering in \\ref{res} and discuss cosmological application of the multiphase strengths of clusters in Section \\ref{omega}. In Section \\ref{conc} we present our conclusions. ",
        "conclusions": "\\label{conc}  We have studied the ability of X-ray emission line ratios to probe the temperature structure of elliptical galaxies and galaxy clusters using theoretical models, the best available {\\sl ASCA} data, and simulated observations with {\\sl Chandra} and {\\sl ASTRO-E}.  We have emphasized multiphase spectra consisting of either two temperatures or cooling flow models and examined the best means of distinguishing between them. Ratios of lines of K shell transitions have been emphasized over the Fe L lines because of possible remaining inaccuracies in the available plasma codes.  From analysis of isothermal MEKAL models computed with energy resolution 2 eV (FWHM) appropriate for the proposed {\\sl Constellation-X} mission, we found that the most useful ratios of K shell transitions for probing the emission measure distributions of ellipticals and clusters are the H1/He4 (K$\\alpha$) and satellite/He4 ratios of the same element (notation described in Table \\ref{tab.notation}). The H2/H1 and He3/He4 ratios (K$\\beta$/K$\\alpha$) are much less temperature sensitive for temperatures where the transitions are strong. Ratios of similar transitions of different elements (e.g., He4 of S/Fe) are very temperature sensitive but must be used with caution for interpreting temperatures structure since they depend on the relative abundances of the elements. Finally, we also discuss the very strong temperature dependence of a few ratios of Fe L lines between $\\sim 0.7$ - 1.4 keV which are typically the strongest lines and most temperature sensitive ratios of all.  We considered a few simple multiphase models to evaluate how well the various line ratios can constrain the general differential emission measure $\\xi(T)$. To facilitate quantitative comparison of ellipticals and clusters having different $\\xi(T)$ we introduced the ``multiphase strength'' parameter, $\\sigma_\\xi$, which measures the fractional width of $\\xi(T)$: $\\sigma_\\xi=0$ for an isothermal gas and $\\sigma_\\xi\\cong 1$ for a constant emission measure distributed over a large temperature range.  Of the simple multiphase models examined (two temperature, constant emission measure, gaussian emission measure, cooling flows) it is found that the two-temperature models have the most flexibility; i.e.  for a given value of $\\langle T\\rangle$ (emission measure weighted temperature) the two-temperature models can generally provide the best approximation to a given line ratio.  The flexibility of the two-temperature model is illustrated by its ability to match very closely a model consisting of a cooling flow component and a single temperature component corresponding to ambient gas (Figure \\ref{fig.2tvscf}). Cooling flows have nearly maximal multiphase strength ($\\sigma_\\xi\\sim 0.7$), but their K$\\alpha$ line ratios (aside from O) can be approximated by two temperature models within $\\sim 20$ per cent for elliptical galaxies like NGC 4472 and within $\\sim 50$ per cent for clusters like Centaurus. The best means to distinguish between two-temperature models and cooling flows is to use oxygen lines or Fe L lines whose ratios can differ by a factor of $\\sim 10$ between the models.  We have re-analyzed the {\\sl ASCA} spectra of three of the brightest galaxy clusters to assess the evidence for multiphase gas in their cores: M87 (Virgo), the Centaurus cluster, and the Perseus cluster. K$\\alpha$ emission line blends of Si, S, Ar, Ca, and Fe are detected in each system as is significant Fe K$\\beta$ emission.  The Fe K$\\beta$/K$\\alpha$ ratios are consistent with optically thin plasma models and do not suggest resonance scattering in these systems. Consideration of both the ratios of H-like to He-like K$\\alpha$ lines and the continuum temperatures clearly rules out isothermal gas in each case.  To obtain more detailed constraints we fitted plasma models over 1.6-9 keV where the emission is dominated by these K shell lines and by continuum. In each case the {\\sl ASCA} spectra cannot determine whether the gas emits at only two temperatures or over a continuous range of temperatures as expected for a cooling flow. The metal abundances are near solar for all of the multiphase models. For Centaurus and Perseus the Si/Fe ratios are consistent with the meteoritic solar values. We also find that our multitemperature models fitted over 1.6-9 keV, which do not require us to accurately model excess absorption or the Fe L lines, also do not require any excess emission at higher energies from, e.g, an AGN component.  We examined the ability of future X-ray missions to constrain $\\xi(T)$ in ellipticals and clusters. By simulating observations with the {\\sl Chandra} ACIS-I and the {\\sl ASTRO-E} XRS of cooling flow models of NGC 4472 and Centaurus we found that these future instruments will have the ability to distinguish between two-temperature models and cooling flows. However, the lower resolution {\\sl Chandra} ACIS-I (and also the {\\sl XMM} ccds) must rely on constraints from O and Fe L lines for definitive results whereas the {\\sl ASTRO-E} XRS only requires the K$\\alpha$ lines above $\\sim 2$ keV (e.g. Si and S).  We have discussed the evidence for resonance scattering in light of our results for the K$\\beta$/K$\\alpha$ ratios. We also discuss the multiphase strength, $\\sigma_\\xi$, as an indicator of the dynamical states of clusters and its potential use as a probe of $\\Omega_0$ in conjunction with galaxy cluster morphologies.  The difficulty in distinguishing between different multiphase models, specifically a two-temperature plasma from a cooling flow, has been emphasized in this paper.  Presently, there are only a few cases where line ratios have been measured with sufficient precision to merely indicate that the gas is not isothermal. It is imperative that the temperature structures of ellipticals and clusters be considered when interpreting properties (e.g. abundances) derived from spectral fits in these systems."
    },
    "9804/astro-ph9804129_arXiv.txt": {
        "abstract": "Equivalent widths of \\OII\\ 3727\\,\\AA\\ lines are measured in 375 faint galaxy spectra taken as part of the Caltech Faint Galaxy Redshift Survey centered on the Hubble Deep Field.  The sensitivity of the survey spectra to the \\OII\\ line is computed as a function of magnitude, color and redshift.  The luminosity function of galaxies in the \\OII\\ line and the integrated luminosity density of the Universe in the \\OII\\ line are computed as a function of redshift.  It is found that the luminosity density in the \\OII\\ line was a factor of $\\sim 10$ higher at redshifts $z\\sim 1$ than it is at the present day.  The simplest interpretation is that the star formation rate density of the Universe has declined dramatically since $z\\sim 1$. ",
        "introduction": "\\label{sec:cintro}  The \\OII\\ line at 3727\\,\\AA\\ is actually a pair of atomic transitions for singly ionized oxygen, the $^2D_{3/2}$ to $^4S_{3/2}$ at 3726\\,\\AA\\ and the $^2D_{5/2}$ to $^4S_{3/2}$ at 3729\\,\\AA.  The transitions are forbidden, meaning that there is no electric dipole connection between the initial and final states, so the spontaneous emission rates are small, $1.8\\times10^{-4}$ and $3.6\\times10^{-5}~{\\rm s^{-1}}$ for the 3726 and 3729\\,\\AA\\ transitions respectively (Osterbrock 1989).  For this reason, the \\OII\\ line is usually collisionally excited by free electrons in hot nebulae; temperatures $T\\sim 10^4$~K are needed to excite the 3.3~eV transitions.  If the electron density is very low, collisional excitation is rare, whereas if it is very high, excited atoms are more likely to be deexcited by a subsequent collision than by spontaneous emission, so there are critical electron densities $n_{\\rm c}$ at which the transitions saturate observationally, defined to be the electron densities at which the collisional excitation rates equal the spontaneous emission rates.  The critical densities depend on temperature because the collisional excitation cross sections do, but at typical temperatures they are roughly $1.6\\times10^4$ and $3\\times10^3~{\\rm cm^{-3}}$ for the 3726 and 3729\\,\\AA\\ transitions respectively (Osterbrock 1989).  (The fact that the two critical densities are different means that the 3726/3729 line ratio can be used to measure electron density.)  The conditions of temperature and density required to excite the \\OII\\ 3727\\,\\AA\\ line are met in H\\,II regions, clouds of ionized hydrogen heated by massive, young, luminous stars.  For this reason, the \\OII\\ emission of a galaxy is sensitive to its young stellar population, or recent star formation history.  In the local Universe, the relationship between \\OII\\ luminosity and star formation has been calibrated \\begin{equation} \\frac{L_{\\rm\\OII}}{2\\times10^{33}~{\\rm W}} = \\frac{R}{1~{\\rm M_{\\odot}\\,yr^{-1}}} \\label{eq:calib} \\end{equation} where $R$ is the star formation rate (Kennicutt 1992).  This relationship shows a significant galaxy-to-galaxy scatter.  It depends on galaxy dust content because dust absorbs strongly in the ultraviolet; the stellar initial mass function because the \\OII\\ luminosity is tied only to the massive star population; and metallicity because the luminosity in the optically thin line ought to be proportional to oxygen abundance, which in turn depends on a galaxy's age and star formation history.  In this study the \\OII\\ luminosity function and \\OII\\ luminosity density of the Universe are measured as a function of redshift.  These functions constrain the star formation history of the Universe. Previous studies on the star formation history of the Universe have used the metal abundances in quasar absorption systems (Pei \\& Fall 1995) and broadband luminosity density (Lilly \\etal\\ 1996); both of these studies show a much higher star formation rate at $z\\sim 1$ than at the present epoch.  The star formation rate of the local Universe has been estimated through the luminosity density in the H$\\alpha$ line (Gallego et al 1995), which is a good measure of star formation rate (Kennicutt 1992).  However, at redshifts $z>0.3$ the H$\\alpha$ line is no longer accessible by visual spectroscopy and therefore difficult to measure.  The \\OII\\ line is a less reliable measure of star formation rate (Kennicutt 1992) but it has the great advantage that it can be measured by visual spectroscopy over the interesting redshift range $0.3<z<1.3$ where the star formation rate density is thought to be changing rapidly.  These considerations motivate the present work.  Several recent studies have shown that the incidence of strong \\OII\\ emission among bright galaxies increases with redshift (Cowie et al 1996; Heyl et al 1997; Small et al 1997; Hammer et al 1997).  This trend implies a higher \\OII\\ luminosity density in the past, which has also been previously measured directly (Hammer et al 1997).  In what follows, physical quantities are quoted in SI units, with Hubble constant $H_0=100\\,h~{\\rm km\\,s^{-1}\\,Mpc^{-1}}$, in world model $(\\Omega_M,\\Omega_{\\Lambda})=(0.3,0)$.  The only exception are number densities, which are given in $h^3\\,{\\rm Mpc^{-3}}$.  Fluxes and luminosities are given as flux and luminosity densities per logarithmic frequency interval, \\ie, $\\nu\\,S_{\\nu}$ or $\\lambda\\,S_{\\lambda}$ and $\\nu\\,L_{\\nu}$ or $\\lambda\\,L_{\\lambda}$, in ${\\rm W\\,m^{-2}}$ and ${\\rm W}$.  Luminosities are all-sphere (not per-steradian). ",
        "conclusions": ""
    },
    "9804/astro-ph9804153_arXiv.txt": {
        "abstract": "We investigate the effects of the UV background radiation on galaxy formation by using the semi-analytic model including the photoionization process.  The semi-analytic model is based on Cole et al. and we use almost the same parameters as their `fiducial' model.  We find that the UV background mainly affects the formation of dwarf galaxies when $J_{-21}\\ga 1 (J\\equiv J_{-21}\\times 10^{-21}$ ergs cm$^{-2}$ s$^{-1}$ sr$^{-1}$ Hz$^{-1}, J$ is the intensity of the UV background).  Because of the suppression of star formation, the number density of small objects corresponding to dwarf galaxies decreases compared with the case of no UV radiation when the UV background exists until the present epoch.  On the contrary, the UV radiation hardly affects massive galaxies.  This is because the massive galaxies are formed by mergers of small galaxies, which are formed at high redshift when the effect of the UV background is negligible.  This strongly suggests that it is important to consider the merging histories of the galaxies.  On the other hand, when the UV background vanishes at a low redshift ($z\\sim 2$), the number density of small objects is hardly changed but the colour becomes bluer, compared with the case of no UV radiation, because stars are newly formed after the UV background vanishes.  Moreover, we show the redshift evolutions of the luminosity functions and the colour distributions of galaxies.  Because the effect of the UV background is strong at low redshift, we can discriminate between the types of evolution of the UV background by observing the evolution of the luminosity function and the colour distributions, if the UV intensity is sufficiently strong ($J_{-21}\\ga 1$). ",
        "introduction": "The problem of galaxy formation is one of the most important unresolved problems in astrophysics.  Recently, owing to the surprising development of the technology concerned with telescopes (e.g. the {\\it Hubble Space Telescope} ({\\it HST}) and the Keck telescope), we can see the birth and evolution of galaxies.  Thus we can obtain some significant clues to understanding the galaxy formation processes.  It is known that the intergalactic medium (IGM) is highly ionized through the Gunn--Peterson test (Gunn \\& Peterson 1965).  Ultraviolet (UV) background radiation, which ionizes the IGM, also affects the galaxy formation processes by photoionizing gas clouds.  The UV photons penetrate the gas clouds and heat up the gas, then star formation in the gas cloud is suppressed by the UV radiation. Therefore it is important to take into account the effects of the UV radiation when we consider the galaxy formation processes.  The sources of the UV background radiation are considered to be mainly quasi-stellar objects (QSOs).  From the observations of Ly{$\\alpha$} clouds and QSOs, the intensity of the UV background radiation is indirectly measured by the `proximity effect' (Carswell et al. 1984; Bajtlik, Duncan \\& Ostriker 1988; Bechtold 1994; Kulkarni \\& Fall 1993; Williger et al. 1994; Giallongo et al. 1996; Lu et al. 1996). The evolution of the UV intensity at low redshifts ($z\\la 2$) is approximately proportional to $(1+z)^{\\gamma}$, $\\gamma\\sim 4$ from the evolution of the QSO luminosity function (Pei 1995), while at high redshifts ($z\\ga 2$) the evolution of the UV background is not yet determined.  It has been pointed out that the UV background affects the galaxy formation (e.g. Dekel \\& Rees 1987).  Babul \\& Rees (1992) and Efstathiou (1992) have discussed the suppression of star formation by the UV background in dwarf galaxies by considering the evolution of a single galactic gas cloud.  Taking into account the photoionization process, the mass function of cooled objects has been calculated in a hierarchical clustering scenario by Chiba \\& Nath (1994).  Their result is that the number density of massive galaxies decreases compared with the case of no UV background.  This behavior is explained as follows.  In the hierarchical clustering scenario, large clouds generally collapse later than small clouds.  As the mean density of a cloud is assumed to be about 200 times the cosmological background density at the collapsing epoch (see, e.g., Peebles 1993), large clouds have a lower density than small clouds.  Then the UV photons penetrate deeply into the large clouds and suppress the star formation, which leads to the decrease in the number density of the massive galaxies.  However, mergers of galaxies actually occur in the hierarchical clustering model.  If luminous galaxies are formed by the mergers of faint galaxies that could collapse at high redshift with little effects of the UV radiation, the number of such luminous galaxies are probably less affected by the UV background.  Therefore, it is important to consider the merging history of galaxies in order to clarify the effects of the UV radiation.  The semi-analytic approach to the galaxy formation, which includes the merging history of dark haloes and galaxies, has been developed recently (e.g. Kauffmann, White \\& Guiderdoni 1993; Cole et al. 1994; Roukema et al. 1997; Somerville \\& Primack 1998).  It has already been confirmed that the averaged stellar age of each galaxy has only a weak correlation with the luminosity of each galaxy by the semi-analytic model (Kauffmann \\& Charlot 1998).  This result would be inconsistent with the shape of the cold dark matter (CDM) mass spectrum unless the mergers occur.  This fact also motivates us to analyse the effects of the UV background on galaxy formation via the semi-analytic approach.  In this paper, we investigate the effects of the UV background radiation with the semi-analytic model, which is an extension of Cole et al. (1994) based on the block model for merging histories of dark haloes (Cole \\& Kaiser 1988; Nagashima \\& Gouda 1997) by including the photoionization process.  We will show how the luminosity functions and the colour distributions of galaxies are changed by the UV background.  In Section 2, we describe our models.  In Section 3, we show the luminosity function and the colour distribution of model galaxies, and their redshift evolution.  Section 4 is devoted to conclusions and discussion. ",
        "conclusions": "We investigate the effects of the UV background radiation on galaxy formation with the semi-analytic model, which is an extension of Cole et al. (1994) by including the photoionization process.  We find that the UV background radiation mainly affects low-mass objects corresponding to dwarf galaxies.  If the UV radiation field exists until the present epoch (case 2), the number of dwarf galaxies at present becomes small.  In hierarchical clustering scenarios, dwarf-scale galaxies are statistically formed at high redshifts, and such galaxies are merged into luminous objects. Therefore, dwarf galaxies at present are formed at low redshifts. Thus, when the effect of the UV background is strong, the star formation in such dwarf galaxies is suppressed by the UV radiation. On the other hand, luminous galaxies are less affected, because such luminous galaxies are formed in high overdensity regions and formed at high redshift owing to the high overdensity environments when the effect of the UV radiation is negligible.  This represents that it is very important to consider mergers of galaxies.  When the UV radiation vanishes at $z=2$ (case 3), the number of dwarf galaxies hardly changes, but the colour becomes bluer because baryonic gas in clouds recombines and begins to form stars after the UV radiation field vanishes.  These blue galaxies might be related to the `{\\it faint blue galaxies}' (e.g. Ellis 1997).  More realistically, in the case that the UV background evolves as $(1+z)^{\\gamma}$, where $\\gamma=4$ for $z\\leq 2$ and $\\gamma=-1$ for $z\\geq 2$ (case 4), the slope of the luminosity functions in the faint end becomes gentler and closer to the observed one, compared with that in the case of no UV background.  The reason why the slope changes is considered as follows.  When the UV intensity declines after $z=2$, large clouds become free from the effects of the UV background at relatively high redshifts, while small objects are affected until low redshifts.  Therefore, the number of galaxies with intermediate magnitudes ($M_{B}-5\\log(h)\\la -18$) increases, compared with the case when the UV background exists with a constant intensity until the present epoch.  Thus the effect of the UV background remains only in small galaxies.  Similarly, the colour distribution becomes very broad because the epoch at which stars can be formed in the decreasing UV radiation field differs according to the mass of the haloes.  We also show the redshift evolutions of the luminosity functions and the colour distributions.  If the UV intensity is sufficiently strong ($J_{-21}\\ga 1$), the difference between the evolutions of the UV background clearly appears in the difference between the evolutions of the luminosity function.  Moreover, because the effect of the photoionization depends on the mass of galaxies, we obtain much information about the properties of the UV background from the colour distributions for different classes of absolute magnitudes of galaxies.  Although our knowledge of processes of the galaxy formation and nature of the UV background is far from complete at present, it is certain that, when the observations of evolutions of luminosity function and colour distributions are realized, as well as the observations of the evolution of Ly$\\alpha$ clouds, in the near future, they will play a significant role in investigating the evolution of the UV background and its effects on galaxy formation.  Thus we believe that it is a powerful tool for investigating the evolution of the UV background to observe the redshift evolutions of the luminosity function and the colour distribution of galaxies.  It should be noted that our conclusion is contrary to the conclusion of Chiba \\& Nath (1994).  They have reported that the number of massive galaxies becomes small in the UV radiation field.  This difference results from considering the mergers of galaxies.  We emphasize that we must construct a galaxy formation model including the mergers of galaxies, because the epoch of the formation of luminous galaxies is not the same as the epoch of the formation of stars in the luminous galaxies.  At the present moment, we do not know the exact criterion for forming stars with the UV radiation.  In order to study the galaxy formations more quantitatively, we need to investigate the criterion for forming stars when the UV background exists in detail, and also the dependence on the merging history model and the other parameters in the model. We will also need to include the chemical evolution process in order to compare our results with the observations in the near future."
    },
    "9804/astro-ph9804015_arXiv.txt": {
        "abstract": "The abundance of rich clusters is a strong constraint on the mass power spectrum. The current constraint can be expressed in the form $\\sigma_8 \\Omega_m^{\\gamma} = 0.5 \\pm 0.1$ where $\\sigma_8$ is the $rms$ mass fluctuation on 8~$h^{-1}$~Mpc scales, $\\Omega_m$ is the ratio of matter density to the critical density, and $\\gamma$ is model-dependent. In this paper, we determine a general expression for $\\gamma$ that applies to any models with a mixture of cold dark matter plus cosmological constant or quintessence (a time-evolving, spatially-inhomogeneous component with negative pressure) including dependence on the spectral index $n$, the Hubble constant $h$, and the equation-of-state of the quintessence component $w$. The cluster constraint  is combined with  COBE  measurements to identify a spectrum of best-fitting models. The constraint from the evolution of rich clusters  is also discussed. ",
        "introduction": "Rich clusters can be used to constrain cosmological models of large-scale structure formation.  Rich clusters are the largest virialized objects in the universe and, hence, their abundance and evolution can be simply related to the linear mass power spectrum, $P(k)$.  Their x-ray temperature can be used to infer the cluster mass. Then, Press-Schechter (1974) theory can be used to  relate the observed cluster abundance as a function of mass to the $rms$ fluctuation on 8~$h^{-1}$~Mpc scales, $\\sigma_8$.  The result is a constraint on a combination of parameters, $\\sigma_8 \\Omega_m^{\\gamma}$, where $\\gamma$ is a function of model parameters.  In this paper, we obtain a general expression for $\\gamma$ which applies to a wide range of models including the standard cold dark matter (sCDM) model, models with a mixture of cosmological constant ($\\Lambda$) and cold dark matter ($\\Lambda$CDM), and QCDM models with a mixture of cold dark matter and quintessence, a dynamical, time-evolving, spatially inhomogeneous component with negative pressure. An example of  quintessence would be a scalar field ($Q$) rolling down a potential $V(Q)$. It has been shown that a spectrum of $\\Lambda$CDM and QCDM models satisfy all current observational constraints (Wang et al. 1998). Our expression for $\\gamma$ includes dependence on the spectral index $n$, the Hubble constant $h$ (in units of 100 km sec$^{-1}$ Mpc$^{-1}$), and the equation-of-state of the quintessence component $w$.  In Section 2, we discuss our derivation of the cluster abundance constraint, which follows earlier derivations to some degree but includes some new features when applied to quintessence. Some may wish to proceed directly to the resulting constraint, given in Section 3. Our result appears to differ slightly  from some earlier works, but we explain in this section the reasons for  those differences. Our purpose in determining $\\gamma$ is to transform observations into a powerful constraint on models. In Section 4, we discuss the further constraint derived from studying the evolution of cluster abundance from redshift $z=0$ to $z \\sim 1$. In Section 5, we conclude by  applying the cluster abundance constraint on $\\sigma_8$ in combination with the constraint on  COBE normalization of $P(k)$ to pick out a spectrum of best-fitting $\\Lambda$CDM and QCDM models. ",
        "conclusions": "The cluster abundance and evolution constraints, when combined with future measurements of the cosmic microwave background, may be an  effective means of discriminating quintessence and $\\Lambda$ models.  The cosmic microwave background (CMB) anisotropy provides a constraint on the mass power spectrum on the horizon scale. For a given model, this constraint from large-scale anisotropy as measured by the COBE-DMR satellite (Smoot et al. 1992; Bennett et al. 1996) can be extrapolated to obtain a limit on $\\sigma_8$ that is completely independent of the cluster abundance constraint.  In Figure~2, we plot the dependence of $\\sigma_8$ on $\\Omega_Q$.  For each $w$, a different curve is shown.  Along each curve is highlighted the range of $\\sigma_8-\\Omega_Q$ consistent with  the cluster abundance constraint derived  in this paper. Hence, the best-fit models are those  near the middle of the highlighted regions.  These are the same models used as examples  in Figure~1.  Near future satellite experiments, such as the NASA Microwave Anisotropy Probe (MAP) and the ESA Planck mission, will greatly improve upon COBE by determining the temperature anisotropy power spectrum to extremely high precision from  large to small angular scales.  Even a full-sky, cosmic variance limited measurement of the CMB anisotropy, though, may  not be sufficient to discriminate $\\Lambda$CDM from QCDM models.  There is a degeneracy in parameter space such that, for any given $\\Lambda$CDM models, there is a continuous family of QCDM models which predicts the same CMB power spectrum (Huey et al. 1998; White 1998) It is possible that the data points to a QCDM model, say, which lies outside this degenerate set of models, {\\it e.g.}, a model with rapidly varying $w$. However, if the data  points to the degenerate set of models, then it is critically important to find a method of discriminating models further.  Not only does degeneracy mean that $\\Lambda$ cannot be distinguished from quintessence, but also that large uncertainty in $\\Omega_m$ and $h$. Here we wish to illustrate how  cluster abundance may play an important role.   An example of a  ``degeneracy  curve\" is shown in Figure~3. Given a value of $h$ for any one point along the curve, values of $h$ can be chosen for other points along the curve such that the models are all indistinguishable from CMB measurements. The near-future satellites are capable of limiting parameter space to a single degeneracy curve. The CMB anisotropy also narrowly constrains $n_s$, $\\Omega_b h^2$, and $\\Omega_m h^2$. However, even when combined with the cluster abundance constraint and other cosmological constraints from the age, Hubble constant, baryon fraction, Lyman-$\\alpha$ opacity, deceleration parameter and mass power spectrum, a substantial degeneracy can remain. Figure~3 includes a shaded region which exemplifies the range which these models might allow, based on current measurements (Wang et al. 1998; Huey et al. 1998). Because of the uncertainty in cluster abundance at $z=0$ and other cosmic parameters, the overlap between the degeneracy curve and the shaded region allows a wide range of $\\Omega_m$, $h$ and $w$.  Cluster evolution offers a promising approach for breaking the degeneracy. Figure~4 illustrates the variation of $A(M_{1.5})$ as a function of $w$ for models along the degeneracy curve and inside the shaded region of Figure~3. The variation in $A(M_{1.5})$ is nearly 2, corresponding to nearly two orders of magnitude variation in abundance at redshift $z=0.5$. The range of $A(M_{1.5})$  is between -3.5 and -5.5 in this case, but this could be shifted upward or downward by adjusting cosmic parameters. The point is that models which are degenerate in terms of CMB anisotropy are spread out in $A(M_{1.5})$. If the measurements can be refined so that $A(M_{1.5})$ is know to better than $\\pm 0.5$, then cluster evolution may play an important role in discriminating between quintessence and  vacuum density and, thereby, determining $\\Omega_m$ and $h$."
    },
    "9804/astro-ph9804237_arXiv.txt": {
        "abstract": "We summarize the status of a ``targeted'' redshift survey aimed at establishing the properties of galaxies and their large scale distribution in the redshift range $2.5 < z < 3.5$. At the time of this writing, we have obtained spectra of more than 400 galaxies in in this redshift range, all identified using the ``Lyman break'' color--selection technique.  We present some of the first results on the general clustering properties of the Lyman break galaxies. The galaxies are very strongly clustered, with co-moving correlation length similar to present--day galaxies, and they are evidently strongly biased relative to the mass distribution at these early epochs, which is consistent with hierarchical galaxy formation models if Lyman break galaxies trace the most massive halos at $z \\sim 3$. Prospects for large surveys for galaxies beyond $z \\sim 4$ are discussed. ",
        "introduction": "At the time of this writing (December 1997), it has been slightly more than 2 years since the ``Lyman break'' technique (discussed extensively elsewhere, e.g. Steidel \\& Hamilton 1993, Steidel, Pettini, \\& Hamilton 1995, Madau  et al. 1996, etc.), for isolating large numbers of high redshift ($z \\sim 3$) galaxies was demonstrated to work efficiently, based on confirming spectra obtained at the W.M. Keck Observatory (Steidel et al. 1996). Galaxies at $z\\sim 3$ or greater are no longer terribly surprising, as now the highest redshift galaxy has $z=4.92$ (Franx et al 1997) and it probably won't be very long before the $z=5$ barrier is surpassed. However, as exciting as finding distant galaxies can be, the greatest gains for our understanding of the galaxy formation process will come from the feasibility of assembling large samples of galaxies at previously inaccessible redshifts, and to study their properties and distribution with a high level of statistical significance. In view of this, we have spent the ensuing 2 years compiling a large sample of $z \\sim 3$ Lyman break galaxies, selected in a consistent manner over a number of relatively large fields. Our aim in carrying out this survey is to obtain $\\sim 100+$ redshifts per field in each of 5--6 fields, each of size 160 square arc minutes or greater. We anticipate that the present survey will reach completion by the end of 1998.  Our survey for Lyman break galaxies continues to be based upon deep ground--based images obtained in a custom broad--band filter system, $U_nG{\\cal R}$. For reliable identification of high redshift galaxy candidates, we have found that long integrations on 4m--class telescopes are required to reach adequate depth and S/N; most of our images have been obtained at the prime focus of the Palomar 5m telescope, and a high quality data set for a 9\\arcmin\\ by 9 \\arcmin\\ field of view requires about 12 hours of total integration time (7 hours in the $U_n$ band alone), or about two good nights of observing. The identification of candidate $z \\sim 3$ galaxies is based upon a combination of the modeling of the spectral energy distributions of high redshift star--forming galaxies, including the effects of the opacity of the galaxy ISM and the intergalactic medium (cf. Madau 1995), and the results of our spectroscopy, which have resulted in a somewhat modified color--selection window as compared to the purely theoretical one which was used originally (cf. Steidel, Pettini, \\& Hamilton 1995). Figure 1 shows a two--color diagram from one pointing of the Palomar 5m, indicating the region from which we select our samples of Lyman break galaxy candidates. Approximately 5\\% of all galaxies to ${\\cal R}=25.5$ are candidate high redshift ($2.5 < z < 3.5$) galaxies. \\begin{figure} \\plotone{fig1.eps} \\caption{An example of a two--color diagram used to select Lyman break galaxies for follow--up spectroscopy. The shaded region corresponds to the region from which candidate Lyman break galaxies are selected. Of the $\\sim$3300 objects in the field to ${\\cal R}=25.5$, approximately 170 of them satisfy our adopted color criteria. Triangles represent objects for which the $U_n-G$ color is a lower limit. } \\end{figure} \\begin{figure} \\plotfiddle{fig2.eps}{2.5in}{0}{40}{40}{-120}{-80} \\caption{Redshift histogram of spectroscopically confirmed Lyman break galaxies in our $z \\sim 3$ survey. The median redshift is $\\langle z \\rangle = 3.02$. The redshift cutoff near $z \\sim 3.4$ is a result of our color selection criterion in $G-{\\cal R}$, while the lower redshift cutoff is imposed by the requirement that a significant ``break'' must be present between the $U_n$ and $G$ passbands (cf. Fig. 1). } \\end{figure} ",
        "conclusions": ""
    },
    "9804/astro-ph9804001_arXiv.txt": {
        "abstract": "New Johnson-Cousins $UBVRI$ photometry of giants in globular clusters is combined with $JHK$ photometry on the CIT system to produce color sequences for giants from the globular clusters M3, M5, M13, and M92. $UBVRI$ data are also presented for giants in the metal-rich open cluster NGC 6791. These data fill a gap in the literature, especially for the $R$ \\& $I$ bands.  We provide the empirical relations between broad band colors for various [Fe/H] values for metal-poor giants. The color sequences for $U-B$ and $B-V$ show clear separations for different [Fe/H] values. We also find weak, though unexpected, metallicity dependences of $V-R$, $V-I$, and $J-K$ colors. $H-K$ is metal-insensitive. The above colors are plotted as a function of $V-K$, and a literature $(V-K) - T_{\\rm eff}$ relation is given. ",
        "introduction": "Population synthesis models and stellar evolutionary isochrones are the tools which give age estimates for galaxies and star clusters, thus providing a constraint on the age of the universe. The models, however, are dependent upon the transformations from effective temperature ($T_{\\rm eff}$), surface gravity, and abundance to observed colors and magnitudes. In the case of metal-poor giants, these transformations remain inconclusive, despite considerable efforts in the past.  In 1966, Johnson\\footnote{Johnson, H. L. 1966, \\araa, 4, 193}\\markcite{j66} produced a remarkable set of tables that gives color-color sequences in $UBVRIJKLMN$ passbands for local dwarfs, giants, and supergiants (of nearly solar metallicity, since metal-poor stars are rare in the disk), based on his and his collaborators' extensive photometry efforts. He calculated the absolute flux in each passband to obtain bolometric corrections and effective temperatures. For giants, this temperature scale has been subsequently revised (Dyck et al. 1996; DiBenedetto 1993; Ridgway et al. 1980)\\markcite{dbvr96,d93,rjww80} but the color-color sequences are still used in many contexts, including assigning colors to theoretical isochrones and computing colors for integrated light population models.  Since 1966, new passbands have become standard for broadband photometry. These are the Cousins (1980a,b)\\markcite{c80a,c80b} $R$ and $I$ filters at roughly 0.68$\\mu$m and 0.79$\\mu$m, and the CIT (Elias et al. 1982; ``CIT'' stands for California Institute of Technology)\\markcite{efmn82} $H$ filter at 1.6$\\mu$m. With these new filters, it is clear that more comprehensive color-color tables are necessary.  One powerful way to generate a complete color-color table is to compute theoretical line-blanketed stellar spectra and convolve them with filter response functions to synthesize colors (e.g. Buser \\& Kurucz 1979; Bell \\& Gustafsson 1978, 1989; Kurucz 1992).\\markcite{bk1979,bg78,bg89,k92} The only major detriment to this approach is that, owing to the complexity of the problems of opacity and convection, the theoretical fluxes match imperfectly with real stellar fluxes. Systematic color drifts can be seen when theoretical results are compared to an empirical color-color table (cf. Worthey 1994).\\markcite{w94}  Empirical tables are therefore of primary interest, for use by themselves or for use as a check of the theoretical color calibrations. The most complete empirical calibration attempted to date was done by Green (1988)\\markcite{g88} for use in the Revised Yale Isochrones (Green et al. 1987).\\markcite{RYI} For Green's work, solar neighborhood $UBVRI$ photometry was assembled, and a temperature scale was attached via the $R-I$ color. The extension to different metallicities was accomplished by using theoretical color results. The Green calibration is imperfect; while problems with fitting isochrones (Corbally 1996)\\markcite{c96} may be due to the temperature of the evolutionary tracks rather than the color calibration, the Green colors differ from other literature calibrations (Worthey 1994).\\markcite{w94} A problem with bolometric corrections (Tiede, Frogel, \\& Terndrup 1995; Mould 1992)\\markcite{tft95,m92} has also been noticed, but the color calibration has not yet been redone.  As part of an effort to redo this color calibration, a literature search conducted by Worthey and Fisher (1996)\\markcite{wf96} discovered a paucity of observations in the $R$ and $I$ passbands for metal-poor giants. New observations at the Michigan-Dartmouth-MIT (MDM) Observatory were proposed to help fill this $RI$ gap and to provide empirical relations between the various colors for different metallicities. We selected globular clusters of known metallicity, with previous $BV$ (and sometimes $U$) photometry, and specifically stars observed in $JHK$ by Frogel et al. (1983)\\markcite{frogel83} and Cohen et al. (1978)\\markcite{cohen78}. In order to extend the metallicity range of our data to higher [Fe/H] values, NGC 6791 was added to the list. Unreduced $JHK$ photometry exists and, we expect, will soon be available to complete the color sequence for this cluster. The approach in this paper is to concentrate {\\it only} on the Frogel et al. (1983) and Cohen et al. (1978)\\markcite{cohen78} stars in the globular clusters and on NGC 6791 giants previously studied by Garnavich et al. (1994).\\markcite{gvzh94} In this way, we reduce a daunting array of data to a manageable size while retaining full ability to construct $UBVRIJHK$ color sequences from the results.  We describe the observations and data reduction in the next section. In section 3.1, we give our photometry results and, in section 3.2, compare them with data found in the literature. Section 3.3 contains a discussion of the empirical color-color relations for the metal-poor giants. Section 4 contains concluding remarks. ",
        "conclusions": "We have presented color sequences and analytical color-color fits for giants as a function of [Fe/H]. The well-known metallicity sensitivity of $(U-B)_{0}$ is clearly visible in Fig.~\\ref{ubvk}. Our data underscore this high sensitivity, especially in the higher $T_{\\rm eff}$ regime. A somewhat weaker dependence upon [Fe/H] was found for $(B-V)_{0}$. This color shows an interesting metallicity degeneracy at $(B-V)_{0}\\sim 1.3$. The colors $(V-R)_{0}$, $(V-I)_{0}$, and $(J-K)_{0}$ are much less influenced by metallicity, while $(H-K)_{0}$ is solely a function of $(V-K)_{0}$ (and thus $T_{\\rm eff}$). The dependence of $(J-K)_{0}$ and $(V-I)_{0}$ on [Fe/H] was a surprising result; more data would certainly be useful to confirm our findings. The extremely weak metallicity dependences of $(V-R)_{0}$ and $(V-I)_{0}$ make these viable temperature indicators, with the advantage of only requiring an optical detector as opposed to an optical/IR combination."
    },
    "9804/astro-ph9804147_arXiv.txt": {
        "abstract": "The UV-bright star Barnard 29 in the globular cluster M13 was observed for 5300 seconds with the Berkeley spectrometer on the ORFEUS-SPAS II mission in 1996 November--December. The resulting spectrum extends from the interstellar cutoff at 912 \\AA\\ to $\\sim$ 1200 \\AA\\ at a resolution of $\\sim$ 0.33 \\AA. It shows numerous absorption features, both photospheric and interstellar, but no significant emission other than diffuse emission of local origin.  The Kurucz synthetic stellar spectrum that best fits the data has $T_{eff} = 21,000$~K, $\\log g = 3.0$, and [M/H] = $-2.5$. This effective temperature and surface gravity are consistent with previous results, but the derived metallicity is lower than that of other M13 giants, for which [Fe/H] = $-1.60$.  Using high-resolution synthetic spectra, we determine the photospheric abundances of C, S, and Fe, species unobservable in the optical.  We find $\\log \\epsilon$(C) = $6.15 \\pm 0.10$, $\\log \\epsilon$(S) = $5.34 \\pm 0.50$, and $\\log \\epsilon$(Fe) = $5.30^{+0.22}_{-0.26}$.  Again, the Fe abundance is lower than expected.  This anomaly may reflect selective condensation of metals onto dust grains at the end of the AGB phase, as has been suggested for some cooler post-AGB stars with peculiar Fe abundances.\\\\ ",
        "introduction": "\\renewcommand{\\arraystretch}{1.1}  The ultraviolet-bright stars in globular clusters lie in the upper left corner of the color-magnitude diagram. Brighter than the horizontal branch (HB) and bluer than the asymptotic giant branch (AGB), this class of objects includes evolved HB, post-HB, and post-AGB stars. Fewer than 50 globular cluster UV-bright stars are known; of these, only 13--16 may be classified as post-AGB stars \\cite{de_Boer87}. The brief post-AGB phase (typically $10^4$ years; \\citeANP{S83} \\citeyearNP{S81,S83}) is one of the least well understood phases of stellar evolution, yet the nucleosynthesis, dredge-up, and mass-loss processes that occur during this phase are principally responsible for the enrichment of the interstellar medium (ISM), and for carbon and oxygen in particular \\cite{WST89}. To improve our understanding of the final stages of a star's lifetime, as well as the mechanisms by which it enriches the ISM, we observed the UV-bright star Barnard 29 with the Berkeley spectrograph on the ORFEUS-SPAS II mission.  Barnard 29 is a well-studied UV-bright star in the globular cluster M13 \\cite{SS70,de_Boer85,CDK94}. Stellar and cluster parameters are presented in Table \\ref{params}.  Recently, \\shortciteN{CDK94} used an LTE analysis of high-resolution optical spectra of Barnard 29 to derive its atmospheric parameters and abundances.  The authors found $T_{eff} = 20,000$ K and $\\log g = 3.0$, confirming Barnard 29 as a post-AGB star.  The star's helium abundance is nearly solar.  Metal abundances are about 1 dex below solar, except for carbon, which is deficient by some 2.4 dex, and nitrogen, which shows a much smaller underabundance of 0.7 dex (relative to solar).  The authors were able to set only upper limits on the abundance of three important elements, C, S, and Fe.  We use the far-UV spectrum of Barnard 29 to set independent constraints on the star's atmospheric parameters and to determine these elemental abundances.  \\nocite{GA86,CM79,Djorgovski93,Stark92,Arp65p401,CDK94,KSSSLP97} ",
        "conclusions": "The CNO abundance pattern seen in Barnard 29 is observed in a number of red giants in M13, four carbon-poor planetary nebulae, and all of the high-latitude B-type post-AGB candidates \\shortcite{CDK94}.  According to current stellar evolutionary theory, such abundance ratios are expected of stars that leave the AGB before third dredge up brings significant nuclear-processed material to the surface \\cite{IR83}.  We would therefore expect Barnard 29 to have an iron abundance similar to other M13 giants, for which [Fe/H] = $-1.60$ \\cite{KSSSLP97}. Instead, we find that the star has an iron abundance 2 dex or more below solar.  Low metallicities are seen in other UV-bright stars in globular clusters, but not all: for vZ 1128 in M3, [M/H] = $-3.5 \\pm 1.5$, nearly 2 dex below the cluster mean, while BS in 47 Tuc and UV5 in NGC 1851 have metallicities consistent with the cluster mean (\\shortciteANP{DDF94} \\citeyearNP{DDF94,DDF95}).  For BD$+33$\\arcdeg 2642, a planetary nebula central star in the galactic halo, \\citeN{NHK94} derive an abundance distribution similar to that of Barnard 29: He is near solar, C, N, O, Mg, Si are depleted by about 1 dex, and Fe is depleted by 2 dex relative to the sun.  The authors suggest that the low Fe abundance may be the result of gas-dust fractionation, a process proposed to explain the chemical peculiarities in some cooler post-AGB stars.  In this scenario, metals with a high condensation temperature condense into dust grains and are removed by radiation pressure, while elements with lower condensation temperatures remain.  Most models require a binary system to support a circumstellar disk about the AGB star (for a review, see \\citeNP{TWW93p103}), but a single-star model has been proposed by \\citeN{ML92}.  One of its predictions, however, is that $N$(C) $\\approx$ $N$(O) in the resulting iron-poor photosphere, a condition not seen in Barnard 29.  The abundances presented here are derived under the assumption of LTE. Because Barnard 29 is a low-gravity star, non-LTE effects may be significant.  The strength of such effects and their influence on the derived stellar abundances are difficult to determine {\\it a priori}. Future work will address these issues using appropriate non-LTE model atmospheres."
    },
    "9804/astro-ph9804284_arXiv.txt": {
        "abstract": "We consider a general method of deprojecting two-dimensional images to reconstruct the three dimensional structure of the projected object, assuming axial symmetry.  The method consists of the application of the Fourier Slice Theorem to the general case where the axis of symmetry is not necessarily perpendicular to the line of sight, and is based on an extrapolation of the image Fourier transform into the so-called cone of ignorance. The method is specifically designed for the deprojection of X-ray, Sunyaev-Zeldovich (SZ) and gravitational lensing maps of rich clusters of galaxies. For known values of the Hubble constant $H_0$ and inclination angle, the quality of the projection depends on how exact is the extrapolation in the cone of ignorance. In the case where the axis of symmetry is perpendicular to the line of sight and the image is noise-free, the deprojection is exact.  Given an assumed value of $H_0$, the inclination angle can be found by matching the deprojected structure out of two different images of a given cluster, e.g., SZ and X-ray maps. However, this solution is degenerate with respect to its dependence on the assumed $H_0$, and a third independent image of the given cluster is needed to determine $H_0$ as well. The application of the deprojection algorithm to upcoming SZ, X-ray and weak lensing projected mass images of clusters will serve to determine the structure of rich clusters, the value of $H_0$, and place constraints on the physics of the intra-cluster gas and its relation to the total mass distribution.  The method is demonstrated using a simple analytic model for cluster dark matter and gas distributions, and is shown to provide a stable and unique reconstruction of the cluster 3D structure. ",
        "introduction": "Uncovering the three-dimensional (3D) structure {}from observations of projected quantities is a fundamental problem of astrophysics research. The 3D structure of rich galaxy clusters is particularly interesting, as the cluster geometry impacts on determinations of the Hubble constant {}from X-ray and Sunyaev-Zeldovich (SZ) measurements, cluster mass and baryon fraction determinations, and the underlying galaxy orbit structure.  The most straightforward approach when dealing with projected quantities is to assume a spherical symmetry.  The assumption is often applied within the framework of parametric models, for example the standard $\\beta$-model (e.g., \\cite{cavaliere76}; see \\cite{rephaeli95} for a review), or a non-parametric geometric approach (\\cite{fabian81}). Such simplifications can lead to biases in whatever quantity is being extracted {}from the data, especially when the systems under scrutiny might be biased to a certain geometry such as, for example, being detected since they are preferably elongated along the line of sight (hereafter LOS). This bias can be considerable as was illustrated in the particular case of merging clusters (\\cite{roet97}).  Here, we restrict ourselves to the problem of reconstructing the 3D structure {}from maps of the X-ray emission or the SZ decrement (say), assuming the underlying 3D distribution is axially symmetric. We show that the Fourier Slice Theorem (\\cite{rybicki87}), can be extended to general case of an arbitrary angle and the deprojection can be done by applying a simple approximation. In the case of astronomical images the inclination angle is  unknown.  However, with the availability of SZ, X-ray and lensing maps, we can uniquely determine the cluster inclination angle by comparing the deprojections of the various images. This result enables us to ascertain the $H_0$ dependence in the more general case where an axial symmetry is assumed.  We apply this technique for the reconstruction of the 3D cluster structure to simulated SZ and X-ray images in an analytical cluster model. ",
        "conclusions": "We present a new algorithm for constructing the 3D structure of axisymmetric objects {}from projected images.  The method is shown to be especially suited to the deprojection of images of clusters of galaxies, and is deonstrated to be stable in the presence of noise and quite a `wide' cone of ignorance.  It is shown that given an assumed value of $H_0$ and a model relating the gas density and temperature, the cluster inclination angle can be solved given its SZ and X-ray maps. With another image of the cluster, $H_0$ and the inclination angle can be determined independently.  The application of the method to realistic cluster simulations, and to SZ, X-ray and weak lensing observations, will be presented in forthcoming papers.  \\newpage"
    },
    "9804/gr-qc9804067_arXiv.txt": {
        "abstract": "An exact solution for a nonrotating boson star in $(2+1)$ dimensional gravity with a negative cosmological constant is found. The relations among mass, particle number, and radius of the $(2+1)$ dimensional boson star are shown. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804020_arXiv.txt": {
        "abstract": "We analyse the integrated properties of the Mid-IR emission of a complete, optically selected sample of galaxies in the Virgo cluster observed with the ISOCAM instrument on board the ISO satellite. The ISOCAM data allows us to construct the luminosity distribution at 6.75 and 15 $\\mu$m of galaxies for different morphological classes. These data are used to study the spectral energy distribution of galaxies of different type and luminosity in the wavelength range 2000 \\AA - 100 $\\mu$m. The analysis shows that the Mid-IR emission up to 15 $\\mu$m of optically-selected, normal early-type galaxies (E, S0 and S0a) is dominated by the Rayleigh-Jeans tail of the cold stellar component. The Mid-IR emission of late-type galaxies is instead dominated by the thermal emission from dust. As in the Milky Way, the small dust grains emitting in the Mid-IR have an excess of emission if compared to big grains emitting in the Far-IR. While the Far-IR emission of galaxies increases with the intensity of the interstellar radiation field, their Mid-IR emission is non--linearly related to the UV radiation field. The spectral energy distributions of the target galaxies indicate that there is a linear relationship between the UV radiation field and the Mid-IR emission of galaxies for low or intermediate activities of star formation, while the emission from the hot dust seems to drop for strong UV fields. The Mid-IR colour of late-type galaxies is not related to their activity of star formation.  The properties of the dust emission in the Mid-IR seem more related to the mass than to the morphological type of the target galaxy. Since the activity of star formation is anticorrelated to the mass of galaxies, this reflects a relationship between the emission of dust in the Mid-IR and the UV radiation field: galaxies with the lowest Mid-IR emission for a given UV field are low mass, dwarf galaxies. These observational evidences are easily explained if the carriers of the Unidentified Infrared Bands that dominate the 6.75 $\\mu$m emission are destroyed by the intense UV radiation field of dwarf galaxies, although abundance effects can also play a role. ",
        "introduction": "The bulk of the infrared (IR) emission of the interstellar medium (ISM) of galaxies is due to the thermal emission of the dust heated by the interstellar radiation field. Stellar photons are absorbed by the interstellar dust and re-emitted in the IR from $\\sim$ 4 $\\mu$m to $\\sim$ 1 mm. The dust emitting in the Far-IR ($\\lambda$ $\\ga$ 70 $\\mu$m) is probably composed of big grains (with a linear radius $r$ $\\ga$ 200 \\AA) of graphite and silicate in thermal equilibrium with the UV and optical photons of the interstellar radiation field (D\\'esert \\etal ~\\cite{Desert}; Dwek \\etal ~\\cite{D97}). In the range 10 $\\mu$m $\\la$ $\\lambda$ $\\la$ 70 $\\mu$m the emission is usually dominated by very small, three dimensional grains (VSG) mostly composed of graphite (10 \\AA $\\la$ $r$ $\\la$ 200 \\AA), absorbing mainly in the UV, whose emission is however not negligible also at shorter wavelengths (D\\'esert \\etal ~\\cite{Desert}; Dwek \\etal ~\\cite{D97}). In the range  3 $\\mu$m $\\la$ $\\lambda$ $\\la$ 15 $\\mu$m the emission of the interstellar medium is generally dominated by the Unidentified Infrared Emission Bands (UIB). The carriers of the UIBs and associated continuum, often called Polycyclic Aromatic Hydrocarbons (PAHs) although this assignment is uncertain, are heated stochastically by absorption of single photons and temporarily reach very high temperatures, at which most of the emission occurs (D\\'esert \\etal ~\\cite{Desert}; Dwek \\etal ~\\cite{D97}).  The knowledge of the dust emission properties of the ISM strongly increased in the last years thanks to the large amount of data obtained by the IRAS satellite, which made an all--sky survey in 4 different filters, at 12, 25, 60 and 100 $\\mu$m. More recently, the ISO satellite is providing higher quality data in the spectral range $\\sim$ 3 $\\mu$m $<$ $\\lambda$ $<$ 200 $\\mu$m. In order to analyse the properties of the IR emission of different classes of objects and/or of single galaxies, the ISOCAM and the ISOPHOT consortia constructed a coordinated program of observations of an optically selected, complete sample of late-type galaxies in the Virgo cluster. This sample, which includes a large number of galaxies of different morphological type and luminosity, has been selected to be used as a reference for other studies. Furthermore the sample includes objects in the centre and in the periphery of the cluster in order to study the effects of the environment on the dust emission and on the evolution of galaxies. Because all the selected galaxies are at the same distance, the construction of a reference sample is quite direct.  The selected sample has been observed entirely with CAM and partly with PHOT and LWS. In this paper we report on the integrated properties of the galaxies observed with CAM, thus in the spectral range 3 $\\mu$m $<$ $\\lambda$ $<$ 20 $\\mu$m. Preliminary results have been already published in Boselli \\etal (\\cite{BISO97}, Paper I) and Boselli \\& Lequeux (\\cite{BL97}). The data, the analysis of the relationship between the star formation activity and the Mid-IR emission of the few resolved galaxies and the study of the effects of the environment on the properties of the dust emission of late-type galaxies will be discussed in future papers (Boselli \\etal, in preparation). The purpose of the present paper is to present a reference for other ISOCAM observations. The sample is described in Sect. 2; a brief description of the observations and of the data reduction is given in Sect. 3. In Sect. 4 we analyse the Mid-IR luminosity distribution of the target galaxies, their spectral energy distribution (SED), and the relationships between the Mid- and Far-IR emission of galaxies of different types and luminosities. The discussion and the conclusions are given in Sect. 5. ",
        "conclusions": "The availability of an unbiased, optically--selected and complete sample allows us to statistically draw the average Mid--IR properties of \"normal\" galaxies. We present for the first time the Mid--IR luminosity distribution at 6.75 and 15 $\\mu$m of a complete, optically-selected sample of late--type galaxies.  The present analysis clearly shows that the Mid--IR emission of ellipticals and lenticulars (morphological range E--S0/a) is dominated by the Rayleigh-Jeans tail of the photospheric emission of the cold stellar component emitting in the near--IR. This is still true for some early--type quiescent spirals. Figure 3 however does not exclude that some of these early--type objects have dust emitting in the Mid--IR. Madden \\etal (\\cite{M97}), by analysing a large sample of early--type galaxies observed with ISOCAM, have in fact shown that many objects characterized by the presence of atomic or molecular gas, or by nuclear star formation, have a Mid--IR emission already dominated by the dust at $\\lambda$ $>$ 6 $\\mu$m (Madden \\etal (\\cite{M97}); Madden et al. in preparation).   In late-type galaxies the Mid-IR emission is generally dominated by dust. The strong relationship between the integrated UV flux at 2000 \\AA ~and the Mid-IR emission (both normalized to the mass of the galaxy) observed by Boselli \\etal (\\cite{BISO97}) indicates that the very small dust grains and the UIB carriers responsible for the emission at 6.75 and 15 $\\mu$m are heated preferentially by hot, massive stars. This is also confirmed by the good morphological correspondance between the peaks of the emission at 6.75 and 15 $\\mu$m and the H$\\alpha$ images of some nearby, bright well--resolved late--type galaxies (Sauvage et al. 1996, Boselli \\& Lequeux 1997). However these dust grains and UIB carriers can also be heated by the visible light of evolved stars if its radiation density is large, as shown for example by the fact that elliptical galaxies with little UV radiation or the bulge of M 31 also exhibit 6.75 and 15 $\\mu$m dust emission (Madden et al. in preparation; Cesarsky D. et al. in preparation; Uchida et al., 1997).  However in low-luminosity late--type galaxies Boselli et al. (1997a) have shown that the integrated Mid--IR emission does not follow the mean UV radiation field: the Mid-IR emission of these objects does not increase, and may even decrease for high Far-UV fluxes. Figure 7 also shows that for such objects the ratio of the Mid-IR flux to the Far--IR flux decreases at low K' luminosities or small total masses (and strong Far--UV radiation fields). All this points to a lack of very small grains and UIB carriers in low-luminosity late--type galaxies, both per unit mass of gas and relative to the big grains that are responsible for the Far--IR emission. There are two possible explanations for this trend.  One can be the effect of metallicity (Sauvage et al. 1990). One expects a lower dust/gas mass ratio for these galaxies since the grains are mostly made of heavy elements, but this does not necessarily imply a lower small grain/large grain ratio. However such a lower ratio could result from a second--order effect: the lower abundance of carbon with respect to oxygen (and probably to silicon) might lower the abundance of carbonaceous grains which are often believed to dominate the Mid--IR emission in our Galaxy, relative to that of silicate grains which are less efficient emitters (see e.g. Dwek et al. 1997). Then the overall Mid--IR emission will decrease with respect to the Far--IR one which is believed to be always dominated by silicates.  The other possible interpretation of the low Mid--IR/Far--IR emission ratio has been proposed by Boulanger et al. (1988) who suggested that in radiation fields larger than $\\sim$ 50 times that in the solar neighborhood up to 80 \\% of the grains emitting in the Mid--IR can be destroyed by UV photons. Such a high radiation field exists in large fractions of the volume of the dwarf galaxies with K' $>$ 10 mag.  With the present data we cannot chose between these two interpretations, which in fact may hold simultaneously. A study of the ISO data base when available will probably allow to solve this problem."
    },
    "9804/astro-ph9804166_arXiv.txt": {
        "abstract": "We discuss the role of tearing instabilities in magnetic reconnection. In three dimensions this instability leads to the formation of strong Alfvenic waves that remove plasma efficiently from the reconnection layer. As a result the instability proceeds at high rates while staying close to the linear regime. Our calculations show that for a resistive fluid the reconnection speed scales as $V_A Re^{-3/10}$, where $V_A$ is the Alfven velocity and $Re$ is the magnetic Reynolds number. In the limit of vanishing resistivity, tearing modes proceed at a non-zero rate, driven by the electron inertia term, giving rise to a reconnection speed $\\sim V_A (c/\\omega_pL_x)^{3/5}$, where $\\omega_p$ is the plasma frequency and $L_x$ is the transverse scale of the reconnection layer. Formally this solves the problem of fast reconnection,  but in practice this reconnection speed is small. ",
        "introduction": "Reconnection of magnetic field lines is a problem that has been hotly debated for the last forty years. Its critical importance stems from the fact that understanding the origin and evolution of large scale magnetic fields is impossible without a knowledge of the mobility and reconnection of magnetic fields. Standard dynamo theories employ the concept of turbulent diffusion to circumvent the problems associated with the high conductivity of astrophysical plasmas (see Parker 1979, Moffat 1978, Krause \\& Radler 1980). Without some sort of enhanced diffusion flux freezing would be an excellent approximation to the motion of the magnetic field in a highly conducting fluid. To change magnetic field topology, to form large scale fields from the small scale loops produced by turbulent motions, one needs to invoke Ohmic diffusion in some way. As this is usually very slow in astrophysical contexts, the turbulent diffusion paradigm appeals to the notion that whenever field lines are properly intermixed, Ohmic dissipation may be enhanced by introducing a very small magnetic correlation scale.  However, this concept has been widely criticized as ill-founded (e.g. Parker 1992, Zweibel 1998). Strong large scale magnetic fields should prevent magnetic fields of opposite polarity from intermixing by turbulent hydrodynamic motions. Both numerical and analytic studies (see \\cite{CV91}, \\cite{KA92}) confirm that the traditional (\\cite{RSS88}) theory of kinematic dynamos is seriously and fundamentally flawed. On the other hand, observations of the solar corona and chromosphere seem to show that reconnection often takes place at speeds of $\\sim 0.1V_A$ (cf. Dere 1996, Innes et al. 1997 and references contained therein).  Evidently at least some astrophysical plasmas can undergo reconnection on short time scales.  Unfortunately, current proposed solutions to this problem are not satisfactory. The widely cited work by Parker (1992) assumes that the galactic dynamo depends on reconnection in the galactic halo, where it is driven by cosmic rays.  This leaves the problem of reconnection in other astrophysical contexts.  Moreover it is far from obvious that reconnection in the galactic halo can provide the basis for an efficient galactic dynamo.  One of us (\\cite{V95a}) has proposed that rapid reconnection of magnetic flux follows from the formation of intense flux tubes in a turbulent plasma. However, this assumes that reconnection is initially rapid enough to allow the formation of such structures in a small number of dynamical time scales and that the plasma has a high $\\beta$ so that the magnetic field can be distributed intermittently. The first assumption may reasonably be regarded as ignoring the question of fast reconnection, since the conditions for flux tube formation are only marginally weaker than the conditions for fast reconnection when the field is not intermittent.  Even granting these assumptions, while flux tubes formation certainly takes place in the solar convection zone and may also be relevant to accretion discs (\\cite{V95b}), it does not seem to provide a universal solution (Lazarian \\& Vishniac 1996). For instance, ambipolar diffusion can infiltrate material into flux tubes and suppress turbulent pumping.  Estimates of the reconnection rate of such flux tubes using the Sweet-Parker reconnection process (\\cite{S58}, \\cite{P57}) do not guarantee reconnection in less than an eddy turn over time under these conditions.  In our earlier paper (Vishniac \\& Lazarian 1998) we studied the role of ambipolar diffusion in reconnection adopting a simple two dimensional Sweet-Parker (\\cite{S58}, \\cite{P57}) geometry. The reconnection speed was shown to be enhanced, but the level of the enhancement was neither sufficient to account for efficient turbulent pumping nor to satisfy the requirements of the galactic dynamo.  In addition, this process is not relevant to reconnection in the solar chromosphere and corona. These considerations motivate our current study of enhanced reconnection.  In this paper we calculate reconnection speeds in the presence of the tearing instability, when the three dimensional structure of the reconnection region is properly accounted for.  The tearing mode instability, which is a particular type of resistive instability,  was quantitatively described by Furth, Killeen, and Rosenbluth (1963).  It is a generic instability which has been cited in many contexts, and in particular, has often been discussed as a means of explaining reconnection rates associated with solar flaring (e.g. Bulanov, Sakai, \\& Syrovatskii 1979, Dere 1996, Glukhov 1996).  One of the problems with the customary treatment of the tearing modes is that in a two dimensional treatment, after a short period of linear growth, they enter a nonlinear stage that depends on the still unclear, but probably slow, evolution of magnetic islands formed during the linear stage (see \\cite{ML84}). In this paper we note that the case of island formation is really singular and that in the generic case, where the opposing field lines are not perfectly anti-parallel, magnetic islands do not form. Instead, strong Alfvenic waves are produced, which carry fluid away from the reconnection zone  so that the instability starts again for new portions of magnetic flux.  The speed of reconnection is determined by the most rapidly growing tearing modes, which reach the end of their linear growth in the time required to eject magnetic flux in the transverse direction.  Slower modes do not have time to develop as the fluid and the flux are carried away.  Consequently, the instability just barely reaches the nonlinear stage, and a linear analysis of the problem is adequate for a qualitative analysis.  A peculiar feature of tearing modes is that they persist as fluid resistivity vanishes. This formally solves the problem of ``fast'' dynamo, i.e. the dynamo in fluid with resistivity tending to zero, but leaves us with the problem of real world astrophysical dynamo as the reconnection rates that we find are small.  In what follows we discuss the physics of the tearing instability (\\S 2), reconnection in resistive fluids (\\S 3), and very highly conducting fluids (\\S 4).  A summary of our results is presented in section 5. ",
        "conclusions": "In this paper we have shown that three dimensional tearing modes in resistive fluids lead to reconnection speeds that scale as the magnetic Reynolds number to the minus three tenths power, compared to the scaling $Re^{-1/2}$ in the Sweet-Parker model. In addition, we have found that unlike standard Sweet-Parker reconnection, reconnection involving tearing modes persists as the fluid resistivity tends to zero while electron inertia drives tearing modes. This formally solves the problem of ``fast'' dynamo, but is of marginal assistance to actual astrophysical dynamos as the tearing reconnection is not fast.  The treatment so far ignores viscosity effects. Electron viscosity can initiate tearing modes on its own. The corresponding term in Ohm's law is proportional to the second derivative of the current. In the case of tearing modes driven by electron inertia this term may become important when sharp current gradients form. The energy deposited is dissipated by viscous dissipation.  In a recent study Glukhov (1996) has shown that viscosity plays an important role when the current sheet is neutral, i.e. there is no magnetic field in the direction of current. We believe that such situations are rather singular and, in general, there will be a component of magnetic field along current. Moreover, a neutral current sheet is likely to be subjected to a strong interchange instability and therefore its response to slower tearing instability is only of academic interest.  In the present paper we do not treat the highly controversial case of ``forced reconnection'' involving tearing modes (see Hassam 1992) and follow the line of reasoning adopted in FKR and later studies (e.g. Kulsrud \\& Hahm 1982).  Finally, we note that our conclusions differ quite dramatically from those of Strauss (1988), who claimed that tearing modes would lead to fast reconnection, i.e. $V_{rec}\\sim V_A$ on the basis of a weakly nonlinear analysis of tearing mode interactions.  However, there is less to this disagreement than would appear at first glance.  We agree that tearing modes will grow to the point of marginal nonlinearity.  We do not agree that this implies fast reconnection.  Instead, global constraints play a dominant role in setting $V_{rec}$, as they do in the usual Sweet-Parker argument without tearing modes.  We conclude that by themselves tearing modes do not look like a panacea for the problems of the astrophysical dynamo. Indeed, although the enhancement of reconnection speeds is substantial in numerical terms, the consequent reconnection rates are not sufficient to support contemporary dynamo theories. Reconnection in the presence of the tearing modes is still too slow.  Is this a crisis? Probably not.  Our present model is still too simple to treat realistic reconnection geometries. In our next paper we will show that the reconnection speeds are substantially enhanced in the presence of MHD turbulence."
    },
    "9804/hep-ph9804453_arXiv.txt": {
        "abstract": "Grand unified theories  may display multiply interacting fields  with  strong coupling dynamics.  This poses two  new problems: (1) what is  the nature of {\\em chaotic reheating} after inflation,  and (2) how is  reheating sensitive to  the {\\em mass spectrum} of these theories ? We answer these questions in two interesting limiting  cases and demonstrate an increased efficiency of reheating  which  strongly enhances non-thermal  topological defect formation,  including monopoles and domain walls. Nevertheless,  the large fluctuations  may resolve this monopole problem via a  modified Dvali-Liu-Vachaspati mechanism in which  non-thermal destabilsation of discrete symmetries occurs at reheating. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9804/astro-ph9804093_arXiv.txt": {
        "abstract": "We have used the Keck 10m telescope to count objects as a function of image size in two high galactic latitude fields covering 1.5 arcmin$^2$ and reaching 50\\% completeness depths of $K$=24 and $J$=24.5 for stellar sources. Our counts extend $\\sim$1 magnitude deeper in $K$ than those of surveys with other telescopes; complement other Keck surveys in the $K$ band that provide counts at comparable or shallower depths but that have not utilized image structure; and extend by several magnitudes the $J$ band counts from brighter surveys using smaller telescopes that cover larger areas. We find the surface-density of objects at $K$=23 to be higher than previously found ($\\sim$500,000 mag$^{-1}$ deg$^{-2}$), but at $K<22$ to be consistent with other surveys. The slope of the $K$ band counts ($d {\\rm log} A / dm = 0.36$) is similar to others near this depth as well as to our own $J$ band counts (0.35). Counts in $J$ and $K$ bands are both in excess of our empirical no-evolution models for an open universe, with the largest excess observed in $J$. The counts are a factor of 2 higher than mild-evolution models at $J$ and $K$ $\\sim$ 23. The slope of the model counts is {\\it insensitive to the assumed geometry} even in the near-infrared primarily because the {\\it model} counts are dominated by low-luminosity ($< 0.1$ L$^*$) objects at modest redshift ($z<1$) with small apparent sizes (half-light radii $\\leq 0.4$ arcsec, i.e. $<4$ h$^{-1}_{50}$ kpc). The observed counts rise most steeply for these smaller objects, which dominate the counts fainter than $K$=22.3 and $J$=23.3. However, {\\it the greatest excess relative to no-evolution models occurs for the apparently larger objects} which have a median $J-K$ of $\\sim 1.5$. At these depths, the size and colors of such objects correspond equally well to luminous ($\\geq0.1$ L$^*$), {\\it blue} galaxies at $1<z<4$, or progressively more diffuse, {\\it blue}, low-luminosity (0.001-0.1 L$^*$) galaxies at $z<1$.  The majority of these sources are too faint for spectroscopic measurement. Based on optical colors, we can rule out that the excess is due to very low luminosity ($< 0.0001$ L$^*$) {\\it red} galaxies at $z<0.25$. We also find a {\\it deficit} of galaxies with red $J-K$ colors corresponding to non-evolving, luminous, early-type (i.e. ``red envelope'') galaxies at $1<z<3$. Even assuming the deficit is due to their appearance as blue galaxies, they could account only for 10-30\\% of the excess of large, blue galaxies. The nature and redshift distribution of excess large and small galaxy populations at $K$=24 and $J$=24.5 remain indeterminate from these data alone. ",
        "introduction": "Near-infrared galaxy counts are considered by some a panacea for the study of galaxy evolution and cosmology. One reason is that compared to counts at optical wavelengths, near-infrared counts are less sensitive to the presently unknown amounts of dust reddening within galaxies, as well as to the uncertain evolution of stellar populations with look-back time. Moreover, in the near-infrared, $k$-corrections are well known to high redshift because they are based on well studied optical photometry of local galaxies (e.g. Bershady 1995). Yet the flatness of the $K$ band counts relative to optical counts has been a puzzle which has so far eluded satisfactory -- or at least an agreed-upon -- explanation (e.g. Cowie \\etal 1990, Gardner \\etal 1993, Gronwall \\& Koo 1995, Djorgovski \\etal 1995, Yoshii \\& Peterson 1995). This is due to the difficulty of interpreting galaxy counts (e.g. Koo 1990, and references therein), which represent a convolution of a detection function, the multi-variate distribution of rest-frame galaxy properties (such as luminosity, color and size), and the cosmological volume. Thus counts alone -- even in the near-infrared -- are unlikely to provide sufficient information to disentangle these various components and discern the true nature of objects at the faintest limits of observation.  The counts are, however, generally derived from measurements of images, which have the potential to yield additional data.  Colors, for example, have been used to infer redshifts beyond the reach of spectroscopy (e.g. Lilly \\etal 1991). Another approach is to exploit image structure at the depths reachable with 4-10m class telescopes (Lilly \\etal 1991, Colless \\etal 1994), but these efforts are still in their infancy. With excellent seeing of 0.3-0.6 arcsec routinely obtainable in the near-infrared from the ground, object size, surface-brightness, concentration, and asymmetry can all be measured and added to counts, magnitudes, and colors for constraining models to the faintest limits reachable from the largest telescopes. For example, between redshifts of $\\sim$0.8 and 3.5, then there is little change in apparent size for a given metric size (e.g. $<$ 20\\% change for 0.1$<$q$_0$$<$0.5). If, in addition, the tight correlation between metric size and luminosity already observed for bright samples of nearby galaxies persists to higher redshift, apparent size and luminosity should be well correlated. In other words, {\\it apparent size can be used to estimate luminosity.}  There are caveats: A significant space-density of low surface-brightness galaxies, if they exist, will broaden the observed size - luminosity relation, particularly if they are more readily detected in deep images (e.g. Ferguson \\& McGaugh 1995).  Evolution via merging may also alter this relation, as may any wavelength dependence of the observed galaxy size.  Nonetheless, faint surveys sample the moderate-redshift universe at very low luminosities inaccessible to brighter surveys. Therefore it is worthwhile to assess whether the faintest sources are predominantly large or small objects. Naively, one might expect the small objects may represent a class of low luminosity objects at low to moderate redshifts, missed by brighter surveys, while the larger objects may be luminous galaxies at high redshift.  In contrast, recent evidence appears to indicate that luminous (L$^*$), high redshift ($z>3$) galaxies are predominantly small (Giavalisco \\etal 1996, Lowenthal \\etal 1997). These measurements, however, have been made in the rest-frame ultra-violet for a select sample of galaxies; the near-infrared image structure for complete, magnitude-limited samples is not known.  In this paper, we consider these extra dimensions of information by presenting very faint, near-infrared counts in the $J$ and $K$ bands as a function of image size from images taken with the Keck 10m telescope. We have taken advantage of the excellent seeing ($\\sim$0.6 arcsec FWHM), large telescope aperture, and low backgrounds to obtain very deep, but small ($\\sim$0.7 arcmin$^2$) images of two high-galactic latitude fields. To assess the nature of the galaxies in our sample, we construct a null hypothesis for the counts and colors as a function of image size at faint magnitudes by adopting a no-evolution model that is based on observations of brighter, multi-color sample (Bershady \\etal 1994). We then determine how the observed counts and size and color distributions deviate from this well-defined prediction. This approach has the advantage that it obviates all but cosmological model assumptions. Hence the measured differences between model predictions and observations can be attributed cleanly to failings of the model or to galaxy evolution without recourse to fine-tuning model-specific parameters.  The organization of the paper is as follows: Section 2 details our field selection and observations, followed by a description of the image processing (Section 3), object detection (Section 4), and photometry and measurements of counts (Section 5). A data table of sizes and magnitudes for individual sources is presented therein. Our analysis of the observed $J$ and $K$ band counts and model predictions are presented in Section 6, while the sizes, colors and evolution are discussed in Section 7. The main findings of our survey are summarized in the concluding Section 8. An appendix contains a description of how counts are derived from images of non-uniform depth ",
        "conclusions": "Our deep Keck near-infrared images have reached surface densities of $\\sim$ 300,000 mag$^{-1}$ deg$^{-2}$ at $J=23.5$ (corresponding to $K=22.75$), which is equivalent to optical counts at depths of $B = 27$ and $I = 26.5$ (Metcalfe {\\it et al.} 1995, Smail {\\it et al.} 1995, Williams {\\it et al.} 1996). Our $K$ band data go somewhat deeper, and reach a surface density of $\\sim$ 500,000 mag$^{-1}$ deg$^{-2}$ at $K=23$. This surface density is higher than any published ground-based $B$ counts, and within 50\\% of the surface densities for the Hubble Deep Field in $V$ and $I$ bands (Williams {\\it et al.} 1996).  A robust result, not dependent on color or size, is that the $K$ band counts do not roll over by $K=22.5$. The same is true for the $J$ band at comparable surface densities.  The excellent seeing conditions have allowed us to use size and color together to identify the dominant galaxy type contributing to the counts. By the faintest magnitudes, the smallest galaxies (i.e. within the bottom 50\\% of the size range for galaxies to these depths) begin to dominate the counts, and have a median $J-K$ of $\\sim$ 1.6. This trend in size and $J-K$ color of the smallest galaxies is qualitatively anticipated from our no-evolution models based on an empirically-determined, local field galaxy luminosity function.  Such galaxies, according to our models, correspond to relatively low luminosity (L$<$0.1L$^*$) galaxies at $z < 1$; their abundance is due to a relatively steep faint end slope for the $K$-band luminosity function, dominated by blue galaxies at low luminosities.  As long as the volume density of galaxies continues to rise at lower luminosities, the counts should continue to rise and be relatively insensitive in {\\it slope} to the cosmological volume. Even in the absence of evolution (though strong evolution is likely at these depths), improved measures of the faint end of the luminosity function must be obtained before galaxy counts can be used as a cosmological probe.  We have also been able to isolate the dominant galaxy type contributing to the count excess, as reckoned with respect to mild or no-evolution models. The most striking result of our deep near-infrared survey to $K$=23 is that there is a substantial excess of apparently {\\it large} galaxies (i.e. within the top 50\\% of the size range for galaxies at these depths), compared to models with no evolution. These galaxies are very blue in $J-K$ if they are at intermediate or high redshifts ($z>0.5$), but relatively red if they are at very low redshift ($z<0.25$). Hence this implies either (i) a ``new'' population of low redshift, low surface-brightness, low-luminosity red galaxies, (ii) an intermediate redshift population of moderately blue, diffuse galaxies with luminosities in the range 0.001-0.1 L$^*$, (iii) a strongly evolving population of galaxies at high redshift observed at L$\\geq$L$^*$, or (iv) some combination of the three. We can rule out the first option based on the blue optical ($R-I$) colors of a random subset of our sample. However, weighing the contributions from the second and third options is not possible without spectroscopy or perhaps detailed morphological information provided from higher resolution images.  We also find a relative paucity of very red galaxies compared to models with no evolution. This deficit is in a region of $J-K$ color and $K$ magnitude corresponding to early-type galaxies brighter than L$^*$ at $z>1$ in the models. If it is assumed that such distant galaxies have evolved to have bluer colors, then they can account for no more than 30\\% of the excess blue galaxies, large and small. The result is insensitive to the assumed value of q$_0$. Further study of the absence of these ``red envelope'' galaxies at faint magnitudes should provide constraints on the epoch of early-type galaxy formation."
    },
    "9804/hep-ph9804429_arXiv.txt": {
        "abstract": "We consider the formation of topological defects after inflation. In order to take into account the effects of the rescattering of fluctuations, we integrate the classical equation that describes the evolution of a complex scalar field on the two-dimensional lattice with a slab symmetry. The growth of fluctuations during preheating is found not to be enough for defect formation, and rather a long stage of the rescattering of fluctuations after preheating is necessary. We conclude that the topological defects are not formed if the breaking scale $\\eta$ is lager than $ \\sim (2 - 3)\\times 10^{16} {\\rm GeV}$. ",
        "introduction": "Inflation \\cite{Guth,Sato} was invented in order to overcome several problems in the standard hot big bang universe, such as the flatness and  horizon problems. Moreover, harmful topological defects created before inflation were diluted away, and are almost absent in the present universe. If some topological defects such as monopoles or domain walls are produced after inflation, they are disastrous since they will soon dominate the energy density of the universe. Therefore, the reheating temperature cannot be as high as the grand unified theory (GUT) scale so as not to produce the GUT monopoles.  It was recently recognized that topological defects may be formed during preheating \\cite{KLS2,Tkachev,KK2}. Preheating is the very beginning of the reheating process, and occurs because of the parametric resonant effects \\cite{KLS1,Shtanov,Boyan1,Yoshi}. Its essence is an {\\it induced} effect in the sense that the presence of the produced particles stimulates further decay of the coherently oscillating scalar field into those particles \\cite{KK1}. This phenomenon is thus peculiar to bosonic particles that obey the Bose-Einstein statistics.  In the preheating stage, very large non-thermal fluctuations are produced, $\\langle \\delta \\phi^2 \\rangle \\sim c^2 M_p^2$ where $M_p$ is the Planck mass and $c=10^{-2}-10^{-3}$ \\cite{KLS2,Tkachev}. These fluctuations change the shape of the effective potential of the field $\\phi$ to restore its symmetry if the potential $V(\\phi)$ is of spontaneous symmetry-breaking type. Later, when the amplitude of these fluctuations is redshifted away by the cosmic expansion, the symmetry is spontaneously broken and topological defects may be created. Thus, the mechanism for producing the topological defects is somewhat similar to the Kibble mechanism in high temperature theory. However, we showed in Ref.\\cite{KK2} that the amplitude of the fluctuations does not grow larger than that of the homogeneous mode for $V(\\phi)=\\lambda(\\phi^2-\\eta^2)^2$, where $\\lambda$ is the self-coupling constant. It implies that the symmetry restoration does not take place and hence the topological defects are not produced. Therefore, we considered that topological defects are formed through another mechanism which is based on the following two facts. The first one is that there are two minima of the effective potential in the radial direction in the simple U(1) theory. The second is that some fluctuations exist initially at the preheating stage. They are produced during inflation and stretched far beyond the horizon to become classical fields. Therefore, the initial condition for the homogeneous field in some region in the universe is different from that in another region in the universe. The dynamics of each region in the universe is thus different and the final value of the field (the minimum into which the field settles down) is different, which leads to the defect formation.  However, in Ref.~\\cite{KK2}, we used the Hartree approximation so that the interactions between particles with different momenta are not fully included. As we will see later, rescatterings \\cite{KT1,KT2,PR,KLS3,GKLS} are quite important to determine the dynamics of the defect formation as suggested by Refs.~\\cite{Kofman1,Kofman2}. To deal with the effect of the rescattering, we must go on to the lattice simulation \\cite{KT1,KT2,PR}. Therefore, in this paper, we integrated the equations describing the dynamics in a real space on the lattice. There are some difficulties in the simulations on lattice. One of the difficulties is how we can implement the quantum fluctuations of the fields on the lattice. We rely on the idea of Ref.\\cite{KT1} and do not discuss in detail (actually, it is not sensitive to the initial conditions for the fluctuations if the order of magnitude is properly taken \\cite{PR}). The other one is the box size and the lattice size. The lattice size must be small enough for identifying the defects. The box size should not be taken too much smaller than the horizon size at the end of the simulation, otherwise we miss the defects whose density is $O(10)$ per horizon. We also take the proper lattice and box sizes so that we can see the resonance effects. From these requirements, we calculate the evolution of a complex scalar field in a two-dimensional lattice abandoning to calculate it in three dimensions because of the lack of machine memories.  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p1_ho.ps}\\\\[2mm] \\caption[fig-ho]{\\label{fig-ho} Evolution of the real part of the homogeneous field for $\\eta=10^{16} {\\rm GeV} $. } \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p1_fl.ps}\\\\[2mm] \\caption[fig-fl-1]{\\label{fig-fl-1} Evolution of the fluctuations for $\\eta=10^{16} {\\rm GeV}$. } \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p1_spec.ps}\\\\[2mm] \\caption[fig-spec-1]{\\label{fig-spec-1} Spectrum of the fluctuations for $\\eta=10^{16} {\\rm GeV}$. These lines shows $\\tau=800,500,200,130,100$ from top to bottom.} \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{r16_fl.ps}\\\\[2mm] \\caption[fig-fl-2]{\\label{fig-fl-2} Evolution of the fluctuations for $\\eta=3\\times 10^{16} {\\rm GeV}$. } \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{r16_spec.ps}\\\\[2mm] \\caption[fig-spec-2]{\\label{fig-spec-2} Spectrum of fluctuations for $\\eta=3\\times 10^{16} {\\rm GeV}$. These lines shows $\\tau=250,200,150,100$ from top to bottom.} \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{rdef_da.ps}\\\\[2mm] \\caption[fig-r-def]{\\label{fig-r-def} Evolution of defects for $\\eta=2\\times 10^{16} {\\rm GeV}$ on lattice at (a) $\\tau=1000$ and (b) $\\tau=1200$. Dots and circles denote defects and anti-defects, respectively.} \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p15phase_2000.ps}\\\\[2mm] \\caption[fig-phase]{\\label{fig-phase} Phase of the field around the string and anti-string at $x \\sim 1510$ and $y \\sim 3830$ in Fig.~\\ref{fig-p15-def}(d). } \\end{figure}   \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p1_scale.ps}\\\\[2mm] \\caption[fig-p1-scale]{\\label{fig-p1-scale} Changes of the number of the defects for $\\eta=10^{16} {\\rm GeV}$ with various initial conditions of fluctuations. The solid line denotes the average.} \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{p15def_da.ps}\\\\[2mm] \\caption[fig-p15-def]{\\label{fig-p15-def} Evolution of the defects for $\\eta=10^{16} {\\rm GeV}$ on the lattice at (a) $\\tau=1300$, (b) $\\tau=1600$, (c) $\\tau=1800$, and (d) $\\tau=2000$. Dots and circles denote defects and anti-defects, respectively.} \\end{figure}  \\begin{figure}[t] \\centering \\hspace*{-7mm} \\leavevmode\\epsfysize=8cm \\epsfbox{z0_fl.ps}\\\\[2mm] \\caption[fig-fl-3]{\\label{fig-fl-3} Evolution of the fluctuations for $\\eta=10^{16} {\\rm GeV}$ using the Hartree approximation in two dimensions in space. } \\end{figure} ",
        "conclusions": "We have considered the formation of topological defects after inflation. To this end, we have integrated the equation which describes the evolution of a complex scalar field with the potential $V(\\Phi)=\\lambda(|\\Phi|^2-\\eta^2)^2/2$ on the two-dimensional lattice with a slab symmetry in the $z$-direction in order to include both parametric resonance effects and rescatterings of fluctuations. We have found that fluctuations produced during preheating do not lead to the defect formation even if their amplitudes become the same order of magnitude as that of the homogeneous mode at the end of preheating. We also have found that the effects of the rescattering of the fluctuations are essential and a rather long duration of the rescattering is necessary for the topological defects to be formed as suggested by Refs.~\\cite{Kofman1,Kofman2}. Fluctuations in the resonance mode scatter off each other to make fluctuations with higher and lower momenta. They also knock the field off from the zero mode to broaden and smooth the spectrum, which in turn stimulates further decays of the homogeneous mode into all the modes. Therefore, the topological defects are formed in a similar way as the Kibble mechanism. These processes take a lot of time so that it is necessary for the field not to fall into the minimum of the potential too soon after preheating. This constrains the value of the breaking scale. From our calculations, we conclude that the topological defects are not formed if $\\eta \\gtrsim (2 - 3) \\times 10^{16} {\\rm GeV}$.  The constraint on the breaking scale is higher than that of our previous result in Ref.~\\cite{KK2} where the Hartree approximation is adopted. In both cases, the amplitude of fluctuations at the end of preheating is not larger than that of the homogeneous mode (compare Figs.~\\ref{fig-fl-1} with \\ref{fig-fl-3}). There are some differences between the calculations with the Hartree approximation and the lattice simulation during preheating. The fluctuations in the $y$-direction have a similar shape, but the amplitude obtained with the use of the Hatree approximation is a little smaller. The fluctuation in $x$-direction in the calculation with the Hatree approximation does not grow until a very late time, which differs completely from the lattice one. The reason is that most of the rescattering effects are neglected in the Hartree approximation while the rescattering is very efficient for producing fluctuations in both $x$- and $y$-directions in the lattice calculations. Therefore, it is not surprising that the critical breaking scale for the formation of topological defects becomes higher in the lattice simulations.   {\\it Note added}. After submission of this paper, we noticed a paper, Ref.~\\cite{TKKL}, in which numerical calculations were done in three dimensions, and similar results were reported. We agree with the authors of Ref.~\\cite{TKKL} that rescatterings are less effective in two dimensions to some extent. However, there is little difference in the resonance effects. Actually, $\\mu \\simeq 0.141$ in ours [see the footnote below Eq.(\\ref{tau-ph})] while $\\simeq 0.147$ in Ref.~\\cite{TKKL}. Moreover, the maximum value of fluctuations is almost the same. The advantage of two-dimensional simulations is that the box size can be taken as large as the horizon size at the end of calculations while the lattice size can be taken small enough to identify defects. Therefore, we can estimate the average numbers of defects per horizon size."
    },
    "9804/astro-ph9804278_arXiv.txt": {
        "abstract": "It is suggested in this paper that the `bare' strange star might be not bare, and there could be a magnetosphere around it. As a strange star might be an intensely magnetized rotator, the induced unipolar electric field would be large enough to construct a magnetosphere around the strange matter core. This kind of magnetosphere is very similar to that of the rotating magnetized neutron stars discussed by many authors. A magnetosphere will be established very soon through pair production by $\\gamma-B$ or two photon processes after a strange star was born in a supernova explosion. It is emphasized that the fact that the strange star surface can not supply charged particles does not stop the formation of a space charge separated magnetosphere around the bare strange star. An accretion crust is quite difficult to come into being around an isolated strange star. Therefore the observed radio signals of an rotation-powered pulsar may come from a bare strange stars rather than a neutron stars or a strange star with an accretion crust. The idea, that the radio pulsars are the strange stars without crusts, is supported by some observations. For example, the electron-positron annihilation line in the spectrum of the Crab pulsar has been reported (Agrinier et al. 1990; Massaro et al. 1991); the iron emission lines have been observed in many X-ray pulsars but never been reported in $X-$ray emission of radio pulsars. This fact is difficult to be understood if the radio pulsars are the neutron stars where the surface binding energy of iron ions is too low to avoid a ion free-flow from the surface (Neuhauseret et al.1986, 1987). ",
        "introduction": "Since the first radio pulsar, CP 1919, was discovered in November 1967 (Hewish, Bell, et al. 1968), more and more radio pulsars are found, the number of which reaches about 750 (Becker \\& Trumper 1997). These objects are now almost universally believed to be rotating magnetized neutron stars. However, it was suggested that there might be no neutron star, only strange stars (e.g. Alock et al.1986). Hence, one of the very interesting and most fundamental questions is `What is the nature of pulsars?'(i.e. Do the signals of pulsars come from neutron stars or from strange stars?). Unfortunately, the question could still not be answered with certainty even now (Broderick 1998).  Soon after the discovery of pulsars, by removing the possibilities of the pulse periods due to white dwarf pulsation or duo to rapid orbital rotation (see, e.g. a review by Smith 1977), many people widely accept the concept that the pulsars are the neutron stars, which were conceived as theoretically possible stable structures in astrophysics (Landau 1932, Oppenheimer \\& Volkoff 1939). Following this, many authors discussed the inner structure of neutron stars, especially the properties of possible quark phase in the neutron star core (e.g. Wang \\& Lu 1984).  As the hypothesis that strange matter may be the absolute ground state of the strong interaction confined state has been raised (Bodner 1971; Witten 1984), Farhi \\& Jaffe (1984) point out that the energy of strange quark matter is lower than that of matter composed by nucleus for quantum chromodynamical parameters within rather wide range. Hence, strange stars, that might be considered as the ground state of neutron stars, could exist, and the observed pulsars might be strange stars (Alcock et al. 1986, Kettner et al. 1995). Therefore, the question about the nature of pulsars, which seems to have been answered, rises again. Now, there are two kinds of candidates for the pulsar's correspondence object: one is the classical neutron star, and another is the strange star which was proposed about a decade ago.  The key point on this question is to find the difference of the behaviors of strange stars and neutron stars, both observationally and theoretically. As both the strange-star models and neutron-star models for pulsars predict the observed pulsars' mass ($\\sim 1.4 M_\\odot$, from double-pulsars systems) and radius ($\\sim 10^6$ cm, from X-ray bursters), it is hard to differentiate these two type models in observations. The dynamically damping effects (Wang \\& Lu 1984; Dai \\& Lu 1996), the minimum rotation periods (Friedman \\& Olinto 1989), the cooling curves (Benvenuto \\& Vucetich 1991) and the vibratory mode (Broderick et al. 1998) have been discussed in detail in the literature. However, no direct observational clue has yet shown that the pulsars are neutron stars or strange stars.  While, in the terrestrial physics, to search the new state of strong interaction matter, the so-called quark-gluon plasma (QGP), is the primary goal of relativistic heavy-ion laboratory (McLerran 1986, Muller 1995). Many proposed QGP signatures have been put forward in theory and analyzed in experimental data, but the conclusion about the discovery of QGP are ambiguities. More likely, it is suggested in theory that there is a possibility of existing strange hardron cluster (Schaffner et al. 1993) and strangelet (Benvenuto \\& Lugones 1995), however, no experiment has affirmed or disaffirmed the suggestion. This laboratory physics researches should be inspired if the pulsars are distinguished as strange stars rather than neutron stars. Also, the rudimental strangelet in the early universe might have implications of fundamental importance for cosmology (e.g. the dark matters, Witten 1984).  Almost all of the proposed strange star models for pulsars have addressed the case generally contemplated by most authors that the strange star core is surrounded by a normal matter crust (Alcock et. al. 1986; Kettner et al. 1995). The essential features of this core-crust structure is that the normal hadron crust with $\\sim 10^{-5}M_\\odot$ and the strange quark matter core with mass of $\\sim 1.4 M_\\odot$ and radius of $\\sim 10^6$ cm are divided by a $\\sim 200$ fm electric gap. It is believed that a crust can be formed during a supernova explosion (Alcock et al.1986) by accretion. However, as discussed below, after the strange star was born, a magnetosphere could be established soon, and the radiation from the open field lines region would prevent the accretion. Therefore, a crust is difficult to form beyond a newborn strange star.  It is accepted that a strange star without crust will not supply charge particles to develop a rotating space charge separated magnetosphere (Alcock et al. 1986). The reason for this is that the maximum electric field induced by a rotating magnetized dipole, $\\sim 10^{11}$ V cm$^{-1}$, is negligible when being compared with the electric field at the strange matter surface, $\\sim 10^{17}$ V cm$^{-1}$. Pulsar emission mechanisms, which depend on the stellar surface as a source of plasma, will not work if there is a bare quark surface (Alcock et al. 1986). Hence the bare strange stars will not be the observed pulsars. However, there are two points pointed out here: \\begin{itemize} \\item The electric field due to electron distribution near the surface decreases quickly outward, from $\\sim 10^{17}$ V cm$^{-1}$ at the surface to $\\sim 10^{10}$ V cm$^{-1}$ at a very small height of 10$^{-7}$ cm above the surface. Therefore the induced electric field will control most area of the magnetosphere. \\item The magnetosphere can be established through pair production in $\\gamma-B$ or two photon processes. \\end{itemize} So, it is proposed that the bare strange stars could act as the observed radio pulsars.  Five conclusions are obtained in this paper: 1. The bare strange stars may be not bare in fact; a magnetosphere would be settled around the strange stars if the pair production process were taken into account. 2. The magnetosphere of a strange star is very similar to that of a neutron star; the radio pulsars might be the `bare' strange stars rather than the neutron stars if the strange matter hypothesis is correct. 3.Both pulsars with parallel and anti-parallel magnetic axes relative to rotational axes can be observed. 4. The idea, that the radio pulsars are the strange stars, is supported by some observations. 5. It is suspected that the strange stars with normal matter crusts are formed in binary systems; and strange stars with crusts would act as X-ray pulsars or X-ray bursters.  The structure of this paper is as follows. In section 2 the formation of the magnetosphere of the bare strange star is discussed, including the pair production processes, which is very important for the exist of magnetosphere. A comparison of properties between magnetospheres of the trange stars and of the neutron stars is presented in section 3. In section 4, the emission of strange star with a magnetosphere is discussed. Conclusion and discussion are shown in the section 5. ",
        "conclusions": ""
    },
    "9804/astro-ph9804108_arXiv.txt": {
        "abstract": "We have carried out high-precision timing observations of the eclipsing binary \\new\\ in the 3.3 years since its discovery. These data indicate that the orbital period is decreasing at a rate of $\\dot{P}_{b}$ = ($-11\\pm$1)$\\times$10$^{-12}$. If secular, this orbital period derivative implies a decay time for the orbit of only 25\\,Myr which is much shorter than the expected timescale for ablation of the companion. We have also measured the proper motion of the pulsar to be 5$\\pm$3\\,mas\\,yr$^{-1}$.  Assuming the pulsar is at the dispersion-measure distance this implies a very slow transverse velocity $v_t$=(30$\\pm$20)\\,km\\,s$^{-1}$. This combination of low velocity and short orbital period argue against formation of the system in the standard manner and we discuss the implications for its evolutionary history. ",
        "introduction": "The discovery of the millisecond pulsar PSR J2051--0827 (\\cite{sbl+96}) as part of the Parkes survey of the southern sky for low-luminosity and millisecond pulsars (MSPs) (\\cite{mld+96,lml+98}) brought to five the number of pulsars which are eclipsed by winds from low-mass ($\\lta$ 0.1\\,M\\sun) companions. It is only the second such system that is not a member of a globular cluster and has the second-shortest orbital period of any binary-radio pulsar, $P_{b}$=2.38\\,h. The duration of the eclipse, $\\sim$10\\% of the orbital period, and the delays at the eclipse egress and ingress indicate that there is material, probably originating from a wind driven from the companion, well beyond the companion's Roche lobe.  Thus, this system joins both \\orig\\ (\\cite{fst88}) and \\tera\\ (\\cite{lmd+90}) as neutron stars which are in the process of ablating their companions and possibly becoming isolated MSPs (\\cite{rst89a,bv91}).  Pulse times of arrival are extremely sensitive to changes in a pulsar's local environment and along the propagation path. Timing measurements of \\new\\ have already shown that there is structure in the electron column density in the eclipse region (\\cite{sbl+96}) and can provide valuable information on the nature of the eclipses (\\cite{sbl+98}).  The eclipsing behaviour and circularity of the orbit suggest that relativistic effects are not going to have an observable effect on the timing. However, measurements of the changes in the orbital period for both \\orig\\ (\\cite{rt91a,aft94}) and \\tera\\ (\\cite{nt96}) suggest that the effects of either the tidal influence of the companion or the ablation process itself may be evident in the timing data. If the pulsar has a reasonable transverse velocity then, given its distance, $\\sim$1.3\\,kpc, the proper motion would also be apparent. ",
        "conclusions": ""
    },
    "9804/astro-ph9804202_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is now well-accepted that selection effects can lead to significant biases in galaxy surveys, especially leading to detections of only a limited range of surface brightness. Disney\\cite{dis76} and others showed that it is very difficult to detect low surface brightness (LSB) galaxies in typical surveys because they do not register above the background of random sky fluctuations. Similarly, high surface brightness (HSB) galaxies would not show any extension and would be confused with stars. Much work has been done on LSB galaxies\\cite{phil98} but the HSB bias has generally been discounted\\cite{all79}.  In this paper we demonstrate that there is a real bias against HSB galaxies in two specific contexts: in nearby clusters where HSB dwarf galaxies---although resolved---are confused with background giant galaxies and also in field galaxy surveys based on photographic surveys where compact galaxies are unresolved. ",
        "conclusions": ""
    },
    "9804/astro-ph9804034_arXiv.txt": {
        "abstract": "s{ The upcoming satellite missions MAP and Planck will measure the spectrum of fluctuations in the Cosmic Microwave Background with unprecedented accuracy. We discuss the prospect of using these observations to distinguish among proposed models of inflationary cosmology. (Fermilab CONF-98/102-A) } ",
        "introduction": "There has been great interest recently in the topic of parameter reconstruction from the Cosmic Microwave Background (CMB). Most of the emphasis in the literature is on the question of determining cosmological parameters such as the density $\\Omega_0$, the Hubble constant $H_0$, or the value of the cosmological constant $\\Lambda$. The purpose of this talk is to take a different approach and apply the machinery of CMB parameter estimation to the subject of inflationary model building.\\cite{dodelson97} We assume generic features of the universe consistent with inflation (a flat universe, vanishing $\\Lambda$), but allow parameters sensitive to inflation to vary. The goal is to determine how well we will be able to distinguish between popular inflation models using upcoming CMB experiments, in particular the all-sky satellite missions MAP and Planck.  The parameters that are of interest for inflation are the ratio of tensor to scalar fluctuation amplitudes measured at the quadrupole $r \\equiv C^{\\rm tensor}_2 / C^{\\rm scalar}_2$, and the spectral index of scalar fluctuations $n$. Fixed parameters are the density of the universe $\\Omega_0 = 1$, and the present vacuum energy $\\Lambda = 0$. Other parameters are allowed to vary, and we plot the expected sensitivity of NASA's MAP satellite\\cite{MAP} and the ESA's Planck Surveyor\\cite{planck} as ellipses projected onto the $r\\,-\\,n$ plane. ",
        "conclusions": "It is evident from Figure 2 that MAP will allow at least rough distinction between large-field, small-field, and hybrid models. Planck, on the other hand, will allow for the beginnings of real precision work in constraining inflation models. Planck will be capable, for instance, of distinguishing the exponent $p$ in a chaotic inflation model $V\\left(\\phi\\right) \\propto \\phi^p$. David Lyth has presented arguments that an appreciably large $r$ is theoretically disfavored.\\cite{lyth97} We finish by noting that these arguments will be put to observational test within the next few years!"
    }
}